source,target
 If the accretion efficiency is particularly low. taux IS correspondingly closer to the black hole. perhaps at some radius smaller than we can resolve for the inner boundary of our dise in these simulations.," If the accretion efficiency is particularly low, $R_{\rmn{h,max}}$ is correspondingly closer to the black hole, perhaps at some radius smaller than we can resolve for the inner boundary of our disc in these simulations."
 While this will impact upon the time-scale of the variations caused by local mass loss - the rate of mass loss depends on the local Eddington rate. which increases with radius - any variations caused by mass loss at smaller radii will occur on shorter time-scales and not affect the gross long-term behaviour that we find in simulation 3.," While this will impact upon the time-scale of the variations caused by local mass loss - the rate of mass loss depends on the local Eddington rate, which increases with radius - any variations caused by mass loss at smaller radii will occur on shorter time-scales and not affect the gross long-term behaviour that we find in simulation 3."
" Furthermore. the inhomogeneity that starts to appear near the hot/cold boundary rapidly spreads to the whole dise outside 771,44."," Furthermore, the inhomogeneity that starts to appear near the hot/cold boundary rapidly spreads to the whole disc outside $R_{\rmn{h,max}}$."
" The triggering factor for the long time-scale variability is that the Z/j4,44 boundary remains fixed. regardless of its magnitude."," The triggering factor for the long time-scale variability is that the $R_{\rmn{h,max}}$ boundary remains fixed, regardless of its magnitude."
 A related issue is the duration of the initial high central accretion rates. which we find at the beginning of simulations 2 and 3.," A related issue is the duration of the initial high central accretion rates, which we find at the beginning of simulations 2 and 3."
 Whether such an event is present in the observed X-ray light curve in Figure | is debatable: the case rests on whether the initial high state that decays over the first 100 days or so of the outburst is part of a single. coherent event or not.," Whether such an event is present in the observed X-ray light curve in Figure 1 is debatable: the case rests on whether the initial high state that decays over the first 100 days or so of the outburst is part of a single, coherent event or not."
 We can. however. comment upon the mechanism that produces the initial high state in the simulations.," We can, however, comment upon the mechanism that produces the initial high state in the simulations."
 It is caused by the accretion of material that starts off inside J?)winx. and becomes irradiated.," It is caused by the accretion of material that starts off inside $R_{\rmn{h,max}}$, and becomes irradiated."
 It follows that the duration of this initial burst will depend on the viscous time-scale at the position of 7/4. Which in turn depends on the factors discussed above.," It follows that the duration of this initial burst will depend on the viscous time-scale at the position of $R_{\rmn{h,max}}$, which in turn depends on the factors discussed above."
" Here. our choice of accretion efficiency. 1) and AMiaatHu) produce a value of Ai4,44 that is well outside the inner boundary of the simulated disc."," Here, our choice of accretion efficiency, $\eta$ and $\dot M_{\rmn{Edd}}(R_{\rmn{0}})$ produce a value of $R_{\rmn{h,max}}$ that is well outside the inner boundary of the simulated disc."
 If nature conspires to make fusus smaller than this estimate. the duration of the initial burst would be shorter.," If nature conspires to make $R_{\rmn{h,max}}$ smaller than this estimate, the duration of the initial burst would be shorter."
 The postulate that the local wind-loss mechanism must occur at small radii and produce short time-scale variations is indeed supported by the results of simulation 3., The postulate that the local wind-loss mechanism must occur at small radii and produce short time-scale variations is indeed supported by the results of simulation 3.
 The mean wind loss rate in the simulation from radii greater than the radius of the inner boundary (2inc3.75«JOH enp is —2«10PNvr3.," The mean wind loss rate in the simulation from radii greater than the radius of the inner boundary $R_{\rmn{wind}} > 3.75 
\times 10^{11} \rmn{cm}$ ) is $\sim 2 \times 10^{-10} \msol \rmn{yr^{-1}}$."
 This is far less than the estimate of 10. to 10.M;ve1 made by Kotanietal.(2000). to explain the hydrogen column density inferred from absorption line spectra.," This is far less than the estimate of $10^{-6}$ to $10^{-7} 
\msol \rmn{yr^{-1}}$ made by \citet{kot} to explain the hydrogen column density inferred from absorption line spectra."
 It is more likely that this kind of mass loss rate is generated at much smaller radii in the dise. where Apad is much lower.," It is more likely that this kind of mass loss rate is generated at much smaller radii in the disc, where $\dot M_{\rmn{Edd}}$ is much lower."
 A fully realistic model for the viscous processes at work in these dises requires non-ideal magnetohydrodynamies (MHD). as the equivalent value of à is likely to vary with dise radius. vertical displacement from the mid-plane and with time.," A fully realistic model for the viscous processes at work in these discs requires non-ideal magnetohydrodynamics (MHD), as the equivalent value of $\alpha$ is likely to vary with disc radius, vertical displacement from the mid-plane and with time."
 While for numerical reasons the values of a which we have used in this work are a little higher than are usually assumed for these disces. the discrepancy is not marked.," While for numerical reasons the values of $\alpha$ which we have used in this work are a little higher than are usually assumed for these discs, the discrepancy is not marked."
 The entire dise spends almost all of its time in the high viscosity state during an outburst., The entire disc spends almost all of its time in the high viscosity state during an outburst.
 MHD calculations have shown that the magneto-rotational instability can sustain a mean equivalent viscosity parameter of a~0.4 CTout&Pringle 1992)., MHD calculations have shown that the magneto-rotational instability can sustain a mean equivalent viscosity parameter of $\alpha \sim 0.4$ \citep{tou}.
. While we use a.=1 here. the viscosity is dominated by the sound speed. scaling as 7xact.," While we use $\alpha = 1$ here, the viscosity is dominated by the sound speed, scaling as $\nu \propto \alpha c_{\rmn{s}}^2$."
 The slight overestimate in à Is easily compensated by our rather conservative estimate of sound-speed in the hot state (described at the beginning of Section 3)., The slight overestimate in $\alpha$ is easily compensated by our rather conservative estimate of sound-speed in the hot state (described at the beginning of Section 3).
 We close by commenting on other factors that contribute to such incredibly high accretion rates during an outburst., We close by commenting on other factors that contribute to such incredibly high accretion rates during an outburst.
 If. during a long quiescent period the inner regions of the accretion disc are absent through evaporation. the accretion dise may build up a very large reservoir of mass.," If, during a long quiescent period the inner regions of the accretion disc are absent through evaporation, the accretion disc may build up a very large reservoir of mass."
 This could lead to extremely high accretion rates during a subsequent outburst., This could lead to extremely high accretion rates during a subsequent outburst.
 In the case of | 105. however. we expect an unusually high mass transfer rate from the companion star. and this could be a fundamental factor in the ability of such systems to achieve and maintain accretion rates near. and perhaps beyond. the Eddington limit.," In the case of $+$ 105, however, we expect an unusually high mass transfer rate from the companion star, and this could be a fundamental factor in the ability of such systems to achieve and maintain accretion rates near, and perhaps beyond, the Eddington limit."
 Mass-transfer rates of the order ~107L;vr.| are only achievable in low-mass, Mass-transfer rates of the order $\sim 10^{-8} \msol \rmn{yr^{-1}}$ are only achievable in low-mass
Our adopted (My. (V-I)) calibration is shown in Figure 10.,"Our adopted $_V$, (V-I)) calibration is shown in Figure 10."
 We match the observations usine a composite relation. Combining the following three polyuouials: As discussed in previous papers (PMSU2: Reid Gizis. 1997). this tripartite approach is required by the uoticeable steepeniug of tlie main sequeuce at (V-I)e2.85.," We match the observations using a composite relation, combining the following three polynomials: As discussed in previous papers (PMSU2; Reid Gizis, 1997), this tripartite approach is required by the noticeable steepening of the main sequence at $\sim 2.85$."
" Finally. Figure LL plots the (M,. (I-J)) relation."," Finally, Figure 11 plots the $_I$, (I-J)) relation."
 There is clearly au abrupt change in slope at (1-J)7-1.5. aud. we have derived separate mean relations for the brighter aud fainter stars. The main sequeuce is esseutially vertical in region of overlap. with au almost even distribution oL datapoints over the range (1.15«(4—J)1.05. 9.2«Al; 11.2).," There is clearly an abrupt change in slope at $\sim1.5$, and we have derived separate mean relations for the brighter and fainter stars, The main sequence is essentially vertical in region of overlap, with an almost even distribution of datapoints over the range $1.45 < (I-J) < 1.65$, $9.2 < M_I < 11.2$ )."
 Rather than attempt to fit a mean relation. we assign an absolute iuagnitude estimate of A;=10.240.7 lor NLTT stars falling in this colour range.," Rather than attempt to fit a mean relation, we assign an absolute magnitude estimate of $M_I = 10.2\pm0.7$ for NLTT stars falling in this colour range."
 The disk main sequence does uot. unfortunately. present a simple linear relation in diagrams - heuce the necessity for the polyuouial relatious computed in the previous section.," The disk main sequence does not, unfortunately, present a simple linear relation in colour-magnitude diagrams - hence the necessity for the polynomial relations computed in the previous section."
 Before applying those calibrations to derive photometric parallaxes [or the NLTT stars. we briefly cousider both the interpretation of the changiug slope of the iain sequence evideut iu Figures 9. 10 aud 11. aud the implications lor our analysis.," Before applying those calibrations to derive photometric parallaxes for the NLTT stars, we briefly consider both the interpretation of the changing slope of the main sequence evident in Figures 9, 10 and 11, and the implications for our analysis."
(Madan ct 11996: Friagaa Torlevich 1998: Steidel 1999).,(Madau et 1996; Friaçaa Terlevich 1998; Steidel 1999).
 The metallicitics presented here aud by Dietrich et ((2002a) based on the emission lines iu quasars at 2Ll are consistent with previous cluission line studies of 2£zXE quasar sauiples., The metallicities presented here and by Dietrich et (2002a) based on the emission lines in quasars at $z\ga 4$ are consistent with previous emission line studies of $2 \la z \la 4$ quasar samples.
 Tn particular. there is no evidence for a decline in the metallicity from 2 to. >L (see also Dietrich IBbuuaun 2003).," In particular, there is no evidence for a decline in the metallicity from $z\simeq 2$ to $z>4$ (see also Dietrich Hamann 2003)."
 Iun the context of galaxy evolution models. hieher metallicities aud shorter evolution timescales are expected for more massive stellar svstenis (Cuedin Ostriker 1997: Cen Ostriker 1999: I&auffinanu Haehnelt 2000: Nolan et 22001).," In the context of galaxy evolution models, higher metallicities and shorter evolution timescales are expected for more massive stellar systems (Gnedin Ostriker 1997; Cen Ostriker 1999; Kauffmann Haehnelt 2000; Nolan et 2001)."
 The close counection of quasars aud the formation of massive galaxies is supported by the relation of the black hole mass aud the mass of the splieroidal ealaxv component (c.g. CGobhardt et 22000: Ferrarese Merritt 2001).," The close connection of quasars and the formation of massive galaxies is supported by the relation of the black hole mass and the mass of the spheroidal galaxy component (e.g., Gebhardt et 2000; Ferrarese Merritt 2001)."
 For elliptical galaxies. a mass inetallicitv relation has been well known for several decades (Saudage 1972: Faber 1973: Bica et 11988).," For elliptical galaxies, a mass – metallicity relation has been well known for several decades (Sandage 1972; Faber 1973; Bica et 1988)."
 lence. a similar mass netallicity relation can be expected for quasars.," Hence, a similar mass – metallicity relation can be expected for quasars."
 Indeed. there Is sole evidence for a correlation of metallicity and linunosity. be. with the black hole mass of quasars. based on broad emüssiou line studies similar to the analysis we present iu this paper (Tamanun Ferland 1993: Shenuuer Netzer 2002: Warner et 22003).," Indeed, there is some evidence for a correlation of metallicity and luminosity, i.e., with the black hole mass of quasars, based on broad emission line studies similar to the analysis we present in this paper (Hamann Ferland 1993; Shemmer Netzer 2002; Warner et 2003)."
 We investigated rest-frame ultraviolet spectra with moderate spectral resolution of a sample of TO hieh redshift quasars with :c3.5., We investigated rest-frame ultraviolet spectra with moderate spectral resolution of a sample of $70$ high redshift quasars with $z \geq 3.5$.
 We used cuussiou-line fux ratios involving carbon. nitrogen. oxvgen. :ib helimu to estimate the metallicity of the linc-ciuitting eas.," We used emission-line flux ratios involving carbon, nitrogen, oxygen, and helium to estimate the metallicity of the line-emitting gas."
 To trausftorm. the observed line ratios iuto metallicities. we used the results of detailed photoionization calculatious as described by Taman ct ((2002).," To transform the observed line ratios into metallicities, we used the results of detailed photoionization calculations as described by Hamann et (2002)."
 A comparison of the eas chemical composition derived front ciission line ratios involving N11] and Hudicates reasonable consistent estimates of the eas nmetallicitv., A comparison of the gas chemical composition derived from emission line ratios involving ] and indicates reasonable consistent estimates of the gas metallicity.
 The estimates of the chemical abundances based ou ΠΟΠΗ audvi. NvitOvilCie)h and ddif ferby~50 Based ou eight individual cussion liue ratios we estimated an average overall uctallicity for the 70 lieh redshift quasars of roughly Z/Z..=~Ltoh.," The estimates of the chemical abundances based on ] and, $+$ ), and differ by $\sim 50$ Based on eight individual emission line ratios we estimated an average overall metallicity for the 70 high redshift quasars of roughly $Z/Z_\odot \simeq 4 \,{\rm to}\,5$."
 Asstuning an upper limit contribution of scattered Lye cinission of Z30 (Hamann Ixorista 1996). the average metallicity of the BELR eas at high redshifts is still uper-solar. within ~20 estimate above.," Assuming an upper limit contribution of scattered $\alpha$ emission of $\la 30$ (Hamann Korista 1996), the average metallicity of the BELR gas at high redshifts is still super-solar, within $\sim 20$ estimate above."
 Compared to previous studies. we find no evideuce for an evolutionary trend iu quasar iuctallicities froin ;c2 fo +2 ," Compared to previous studies, we find no evidence for an evolutionary trend in quasar metallicities from $z\simeq 2$ to $z\ga 4$."
We analyze the derived. clemental abuudauces witlin the context of models preseuted by IEuuauu Ferland (1993) and Friagaa Terlevich (1998)., We analyze the derived elemental abundances within the context of models presented by Hamann Ferland (1993) and Friaçaa Terlevich (1998).
" With au evolution time scale of approxiuatelv Τι~0.5toOS CC, the epoch of the first intense star formation is estimated to beein as carly as at a redshift of iyc6toδ, he. less than 1 Cer of the age of the universe (IL,z65 bAIMpe 104, =0.3. 04=0.7)."," With an evolution time scale of approximately $\tau _{evol} \sim 0.5 ~{\rm to}~ 0.8$ Gyrs, the epoch of the first intense star formation is estimated to begin as early as at a redshift of $z_f \simeq 6 ~{\rm to}~ 8$, i.e., less than 1 Gyr of the age of the universe $_o \simeq 65$ $^{-1}$ $^{-1}$, $\Omega _M = 0.3$, $\Omega _\Lambda = 0.7$ )."
 We find a weal trend of aL Z/Z.. relation for the high redshift quasars., We find a weak trend of a $L$ – $Z/Z_\odot$ relation for the high redshift quasars.
 Duc to the scatter in nietallicitv and the small rauge of covered huninosity the probability for a correlation bx chance 1s 6 However. the treud is in good agreciment with results obtained for quasars at lower redshift (Ihunanu Ferland 1999) and for composite spectra based on a large quasar sample. which we are currently investigating (Dietrich et 2002b: Warner et 22003).," Due to the scatter in metallicity and the small range of covered luminosity the probability for a correlation by chance is $\sim 6$ However, the trend is in good agreement with results obtained for quasars at lower redshift (Hamann Ferland 1999) and for composite spectra based on a large quasar sample, which we are currently investigating (Dietrich et 2002b; Warner et 2003)."
As a consequence we expect to detect cemission in all hieh density rregions or equivalent in all voung CY clusters.,As a consequence we expect to detect emission in all high density regions or equivalent in all young $UV$ clusters.
 Our analvsis supports a scenario in which the interaction between the XD galaxv NGC 1510 and. the large spiral galaxy. NGC 1512 has triggered star formation activitv in the outskirts of the disk and enhanced the tidal distortion in the aarms., Our analysis supports a scenario in which the interaction between the BCD galaxy NGC 1510 and the large spiral galaxy NGC 1512 has triggered star formation activity in the outskirts of the disk and enhanced the tidal distortion in the arms.
 Phe interaction seems to occur in the north western areas of the system because of the broadening of the aarm and the spread. of the CV -rich star clusters in. this region., The interaction seems to occur in the north western areas of the system because of the broadening of the arm and the spread of the $UV$ -rich star clusters in this region.
 The svstem is probably in the first stages of a minor merger which started 7400 Myr ago., The system is probably in the first stages of a minor merger which started $\sim$ 400 Myr ago.
 Future ssurvevs. such as those planned with the Australian SIVA Pathfinder CASIADI: Johnston et al.," Future surveys, such as those planned with the Australian SKA Pathfinder (ASKAP; Johnston et al."
 2008) will produce similar ccubes and images han obtained here. but over much larger areas.," 2008) will produce similar cubes and images than obtained here, but over much larger areas."
 E.g.. the Xoxosed. shallow ASIAP ssurvey of the sky will reach a sensitivity. of —1 + a an angular resolution. of⋅⋅∕∕ iin a 12-h integration per field.," E.g., the proposed shallow ASKAP survey of the sky will reach a sensitivity of $\sim$ 1 $^{-1}$ at an angular resolution of in a 12-h integration per field."
 Focal plane arrays will provide a very large. instantaneous field-of-view of 525.575.," Focal plane arrays will provide a very large, instantaneous field-of-view of $5\fdg5 
\times 5\fdg5$."
 ‘This means that iiniages similar to those shown in this paper will be obtained for the entire Local Volume., This means that images similar to those shown in this paper will be obtained for the entire Local Volume.
 Furthermore. the correlator bancwidth of 300 Alllz (divided into 16.000 channels) will allow us to σεν the ccontent οἱ ealaxies and their surroundings out to 760.000 (ές = 0.2).," Furthermore, the correlator bandwidth of 300 MHz (divided into 16,000 channels) will allow us to study the content of galaxies and their surroundings out to $\sim$ $z$ = 0.2)."
 In ackition. very deep 20-cmi radio continuum images are obtained for the same area.," In addition, very deep 20-cm radio continuum images are obtained for the same area."
a one-zone model.,a one-zone model.
 The collision rates of CO are from Flower&Launay(1985) for temperatures frou 10 to 250 TN. aud from Mckeeetal.(1982) for 500 to 2000 Is. We asstune CO aud ®CO abundances with respect to Ils of «10 aud 1410© with the observed velocity exadient of — 1 luis | peb of the ring.," The collision rates of CO are from \citet{flow85}
 for temperatures from 10 to 250 K, and from \citet{mckee82}
 for 500 to 2000 K. We assume $^{12}$ CO and $^{13}$ CO abundances with respect to $_{2}$ of $\times10^{-5}$ and $\times10^{-6}$ , with the observed velocity gradient of $\sim$ 1 km $^{-1}$ $^{-1}$ of the ring."
 We determined the velocity eradieut in the Paper I by the PV diagram. and it is consistent in this paper.," We determined the velocity gradient in the Paper I by the PV diagram, and it is consistent in this paper."
 The average ratio of the narrow and broad line clumps are used., The average ratio of the narrow and broad line clumps are used.
 Cluups Nt. 01. D1. D2. DL. D5 are excluded in the average ratio because of their large uncertainty in Ris.," Clumps N4, B1, D1, D2, D4, D5 are excluded in the average ratio because of their large uncertainty in $_{13}$ ."
 Therefore the average Re. aud Ry of the uarrow Lue clamps are 1.00-£0.02 and 9.9042.11. respectively.," Therefore the average $_{32}$ and $_{13}$ of the narrow line clumps are $\pm$ 0.02 and $\pm$ 2.11, respectively."
" The average B3» aud Ry: of the broad Lue clips are (7240.01 and 9.551.456. respectively,"," The average $_{32}$ and $_{13}$ of the broad line clumps are $\pm$ 0.01 and $\pm$ 1.56, respectively."
 With the constraint of the inteusitv ratios within the uncertainty. the estimated temperature aud density. of the narrow line clumps are 2250 [EK aud (L543.5)«105 7.," With the constraint of the intensity ratios within the uncertainty, the estimated temperature and density of the narrow line clumps are $\ge$ 250 K and $(4.5\pm3.5)\times10^{3}$ $^{-3}$."
 The broad line chuups have temperatures of 152-15 IK and deusity of (8.541.5)«10? cu7., The broad line clumps have temperatures of $45\pm15$ K and density of $(8.5\pm1.5)\times10^{2}$ $^{-3}$.
 The predicted brightness temperature (5) is 100 K for the narrow line chuups and —20 Is for the broad line clumps., The predicted brightness temperature $T_{\rm b}$ ) is $\sim$ 100 K for the narrow line clumps and $\sim$ 20 K for the broad line clumps.
" However. it soenis to be inconsistent with the high/low Xj, aud low/high uunuber density iu the broad/narrow line chips if we assune a coustaut scale height for the chimps."," However, it seems to be inconsistent with the high/low $\Sigma_{\rm H_2}$ and low/high number density in the broad/narrow line clumps if we assume a constant scale height for the clumps."
 The solution iav be a sanaller beam filling factor for the narrow line clamps., The solution may be a smaller beam filling factor for the narrow line clumps.
 Iu Figure 9cc. the B3» values have a positive correlation with Xagg.," In Figure \ref{fig-clump-sfr}c c, the $_{32}$ values have a positive correlation with $\rm\Sigma_{SFR}$."
 In Figure 9dd. similar to Mapp. Reo is slightly lower iu the broad line ring chuups aud does not show any svsteniatie pattern in the azinmthal direction.," In Figure \ref{fig-clump-sfr}d d, similar to $\rm\Sigma_{SFR}$, $_{32}$ is slightly lower in the broad line ring clumps and does not show any systematic pattern in the azimuthal direction."
 Figure baa is the intensity weighted isovelocity map of 12000] 2 21)., Figure \ref{fig-mom1}a a is the intensity weighted isovelocity map of $^{12}$ CO(J = 2–1).
 The gas motion iu the ring appears to be donated by circular motion. while it shows clear nou-circular motions in the τςΟΙ = 10) lap as indicated by the S-shape uearly parallel to the dust lanes.," The gas motion in the ring appears to be dominated by circular motion, while it shows clear non-circular motions in the $^{12}$ CO(J = 1–0) map as indicated by the S-shape nearly parallel to the dust lanes."
 As we discussed in Paper I. the non-significant non-circular motion in the COUT = 21) maps is perhaps because the dust lanes are uot as strouelv detected in οἱ) = 21) line. along with the fact that they are closer to the edee of our primary beam. or the non-circular motion is not pronmuneut at the high spatial resolution.," As we discussed in Paper I, the non-significant non-circular motion in the $^{12}$ CO(J = 2–1) maps is perhaps because the dust lanes are not as strongly detected in $^{12}$ CO(J = 2–1) line, along with the fact that they are closer to the edge of our primary beam, or the non-circular motion is not prominent at the high spatial resolution."
 The cincunmnuclear gas is in general in solid body rotation., The circumnuclear gas is in general in solid body rotation.
 The velocity eradieut of the blueshifted part is slightly steeper than the redshifted part., The velocity gradient of the blueshifted part is slightly steeper than the redshifted part.
 We also show the intensity weielitted velocity clispersion map in Figure 5bb. As we inentioned above. the velocity dispersion is larger in the twin-peak region. and lower iu the region away from the twiu-peak region.," We also show the intensity weighted velocity dispersion map in Figure \ref{fig-mom1}b b. As we mentioned above, the velocity dispersion is larger in the twin-peak region, and lower in the region away from the twin-peak region."
 The dynamical center of NCC 1007 was derived by Koloetal.(2003) in thei low resolution 12C0(J = 10) map.," The dynamical center of NGC 1097 was derived by \citet{koh03}
 in their low resolution $^{12}$ CO(J = 1–0) map."
" With our high resolution. !1?CO(J = 21) nap. we expect to deteriune the dvaiiuuical center more accurately,"," With our high resolution $^{12}$ CO(J = 2--1) map, we expect to determine the dynamical center more accurately."
 We use the AIPS taskGAL to determine the dynamical center., We use the AIPS task to determine the dynamical center.
 In the taskGAL. ο = 21) intensitvaweielted velocity map (Figure 5)) isused to fit a rotation curve.," In the task, $^{12}$ CO(J = 2–1) intensity-weighted velocity map (Figure \ref{fig-mom1}) ) isused to fit a rotation curve."
 The deduced kinematic parameters are sununarized in Table 6., The deduced kinematic parameters are summarized in Table 6.
 We use au exponential curve to fit the area witli 177 m radius;, We use an exponential curve to fit the area within $\arcsec$ in radius.
 The observed rotation curve aud the fitted model curve are shown in Fieure 15.., The observed rotation curve and the fitted model curve are shown in Figure \ref{fig-gal}.
 From the fitted parameters. we find that the offset (~ 073) of the dvuiuica center with respect to the position of the AGN is still within a fraction of the svuthesized beam size.," From the fitted parameters, we find that the offset $\sim0\farcs3$ ) of the dynamical center with respect to the position of the AGN is still within a fraction of the synthesized beam size."
 The derived Vy. has a differcuce of ~5 kins ! between @CO(J = 1.0) aud 132C0(J = 21) data. which is less than the velocity resolution of the data.," The derived $V_{\rm sys}$ has a difference of $\sim$ 5 km $^{-1}$ between $^{12}$ CO(J = 1–0) and $^{12}$ CO(J = 2–1) data, which is less than the velocity resolution of the data."
 Upon examiniug the chaunel maps of the ΟΙ = 21) data. we find that the peak of the nuclear emission is aluost coincident with the position of the AGN with an offset of 0733.," Upon examining the channel maps of the $^{12}$ CO(J = 2–1) data, we find that the peak of the nuclear emission is almost coincident with the position of the AGN with an offset of 3."
 We therefore conclude that tle position offset iu the integrated intensity map. as mentioned in Sect. 3.1.," We therefore conclude that the position offset in the integrated intensity map, as mentioned in Sect. \ref{sect-morphology},"
 is due to the asvuuuetric intensity distribution., is due to the asymmetric intensity distribution.
 Iu the low resolution CO maps (Paper LE. IXolnoetal.2003)). NGC 1097 shows bright CO twin-peak structure arising at the intersection of the starburst rine aud the dust lanes.," In the low resolution CO maps (Paper I, \citealt{koh03}) ), NGC 1097 shows bright CO twin-peak structure arising at the intersection of the starburst ring and the dust lanes."
 The 2300 pe resolution CO data show that the barred galaxies usually lave a laree amount of central concentration of molecular gas (Sakamotoetal.1999)., The $\ge$ 300 pc resolution CO data show that the barred galaxies usually have a large amount of central concentration of molecular gas \citep{sakamoto99}.
. IXenuev.etal.(1992) found that im several barred galaxies which host circunuuclear rines (MIOL. NGC 3351. NGC 6951). the central couceutratious. of molecular gas were resolved iuto twin-peak structures when resolution of ~200 pc is attained.," \citet{ken92} found that in several barred galaxies which host circumnuclear rings (M101, NGC 3351, NGC 6951), the central concentrations of molecular gas were resolved into twin-peak structures when resolution of $\sim$ 200 pc is attained."
 A paix of CÓ intensity concentrations are found in these cases. in the cireunnuuclear rine. at the intersection of the rine aud the dust lane.," A pair of CO intensity concentrations are found in these cases, in the circumnuclear ring, at the intersection of the ring and the dust lane."
 Their orientation is almost perpendicular to the major stellar bar., Their orientation is almost perpendicular to the major stellar bar.
 The twin-peak. structure can be attributed to the orbit crowding of iuflowiug eas stream lines., The twin-peak structure can be attributed to the orbit crowding of inflowing gas stream lines.
 The eas flow changes from its original orbit (the so called ο orbit) when it encounters the shocks. which results in a large deflection angle aud migrate to new orbit (the so called cc» orbit).," The gas flow changes from its original orbit (the so called $x_{1}$ orbit) when it encounters the shocks, which results in a large deflection angle and migrate to new orbit (the so called $x_{2}$ orbit)."
 The gas then accumulates in the family of the c» orbits in the shape of a ring or nuclear spirals (Athanassoula19922:Pineretal.1995).," The gas then accumulates in the family of the $x_{2}$ orbits in the shape of a ring or nuclear spirals \citep{atha,piner95}."
. Intense lnassive star formation would follow iu the riug/nuuclear spiral once the gas becomes deuse enough to collapse (Ehucercen1991)., Intense massive star formation would follow in the ring/nuclear spiral once the gas becomes dense enough to collapse \citep{elme94}.
. Tn our 100 pe resolution CO imap. the starburst molecular ring is resolved into individual CALTAs.," In our 100 pc resolution CO map, the starburst molecular ring is resolved into individual GMAs."
 Iu the orbit crowding region. we resolve the twin-peak iuto broad line clamps associated with the curved dust lanes.," In the orbit crowding region, we resolve the twin-peak into broad line clumps associated with the curved dust lanes."
 The narrow line chuups are located away frou the twin-peak and are associated with star formation., The narrow line clumps are located away from the twin-peak and are associated with star formation.
" This kiud of ""spectroscopic components” were also slow ii several twin-peak ealaxics at the intersection of dust lanes and circuuuuclear ring. such as NGC 1365 (Sakamotoetal.2007).. NCC 1151 (Duasetal.2010).. NGC 69016 (Schinnereyretal. 2007).. and NGC 6951 (Isohnoetal1999)."," This kind of “spectroscopic components” were also shown in several twin-peak galaxies at the intersection of dust lanes and circumnuclear ring, such as NGC 1365 \citep{sakamoto07}, NGC 4151 \citep{dumas10}, NGC 6946 \citep{schinnerer07}, , and NGC 6951 \citep{koh99}."
. ILowewer. most of the spectra at these intersections show blended narrow/broad line colponcuts. which is perhaps due to insufficient aueular resolution.," However, most of the spectra at these intersections show blended narrow/broad line components, which is perhaps due to insufficient angular resolution."
 Our observations for the first tine spatially resolved these two compoucuts toward the twiu-peak reeion of NGC 1097., Our observations for the first time spatially resolved these two components toward the twin-peak region of NGC 1097.
 Itis interesting to note that theciemunuuclear ring is nearly circular at ~ 127 inclination. which indicates its intrinsic elliptical shape iu the ealactic plane.," Itis interesting to note that thecircumnuclear ring is nearly circular at $\sim$ $\degr$ inclination, which indicates its intrinsic elliptical shape in the galactic plane."
 The schematic sketch is shown in Figure 17.., The schematic sketch is shown in Figure \ref{fig-model}. .
 T10 loci of dist lanes are invoked to trace the galactic slick wave. aid their shapes are dependent on the paranetersof t1C barred potential.," The loci of dust lanes are invoked to trace the galactic shock wave, and their shapes are dependent on the parametersof the barred potential."
 In the case of NGC 1097. the observed dust lanes resemble the theoretical studies," In the case of NGC 1097, the observed dust lanes resemble the theoretical studies"
", f top)presentsseveralhubsofine spacedaliasperiods.",", top) presents several hubs of fine-spaced alias periods."
"Wecandiscardallhubslongwardsoff 2.5, sincethelongestcontinuousdatasetof 6. 28hclearlydoesnotrep require: er sig))."," We can discard all hubs longwards of $f = 2.5$ , since the longest continuous data set of 6.28 h clearly does not represent 2/3 of an orbit (the triangles in \\ref{ec13phlc_fig}))."
T hehubshortwardso ff 1.5yieldsanellipsoidallightcurveatanorbitalperiodP ~22h., The hub shortwards of $f = 1.5$ yields an ellipsoidal light curve at an orbital period $P \sim 22~\mathrm{h}$.
" At a spectral type of M1 (see below) this would require an evolved secondary, and there is no spectroscopic evidence that would support such a scenario."," At a spectral type of M1 (see below) this would require an evolved secondary, and there is no spectroscopic evidence that would support such a scenario."
 The hub centred at f=2.13 therefore remains as the only possibility., The hub centred at $f = 2.13$ therefore remains as the only possibility.
" We have folded the photometric data with all periods with peak values larger than half the value of the strongest peak, covering a frequency range 1.907—2.381cyc/d."," We have folded the photometric data with all periods with peak values larger than half the value of the strongest peak, covering a frequency range $1.907-2.381~\mathrm{cyc/d}$."
" Based on the criterion of how the data sets of different nights fit together in the phase-folded data, we find that only two periods, Ρι=0.4757d and =0.4695d yield an acceptable light curve."," Based on the criterion of how the data sets of different nights fit together in the phase-folded data, we find that only two periods, $P_1 = 0.4757~\mathrm{d}$ and $P_2 = 
0.4695~\mathrm{d}$ yield an acceptable light curve."
" Since P» is the slightly stronger one of the two, we adopt as photometric period Pp,=0.4695(01)d."," Since $P_2$ is the slightly stronger one of the two, we adopt as photometric period $P_\mathrm{ph} = 0.4695(01)~\mathrm{d}$."
" As a word of caution we remark that our criterion here assumes that each data set represents a part of a stable, identical light curve."," As a word of caution we remark that our criterion here assumes that each data set represents a part of a stable, identical light curve."
" However, the potential presence of star spots or activity on the secondary star could induce a certain variability of the light curve."," However, the potential presence of star spots or activity on the secondary star could induce a certain variability of the light curve."
" This applies to all three targets of this study, but bears special importance for 113349—3237, as here we are dealing with 4 incomplete parts of a light curve within two data sets that are separated by one month."," This applies to all three targets of this study, but bears special importance for 13349–3237, as here we are dealing with 4 incomplete parts of a light curve within two data sets that are separated by one month."
" Somewhat surprisingly, the spectroscopic data do not present a similarly clear variation, and in fact do not appear to reflect the photometric variation at all."," Somewhat surprisingly, the spectroscopic data do not present a similarly clear variation, and in fact do not appear to reflect the photometric variation at all."
" Measuring radial velocities by fitting single Gaussians to the Ha emission line or to a number of absorption lines 145893, CaL16103 and 46122) results in very noisy curves without any clear periodic signal."," Measuring radial velocities by fitting single Gaussians to the $\alpha$ emission line or to a number of absorption lines $\lambda$ 5893, $\lambda$ 6103 and $\lambda$ 6122) results in very noisy curves without any clear periodic signal."
" In a second attempt we measured radial velocities by cross-correlation in the spectral regionA,, which contains a forest of absorption lines from the secondary star due to Mg, Cr and Fe."," In a second attempt we measured radial velocities by cross-correlation in the spectral region, which contains a forest of absorption lines from the secondary star due to Mg, Cr and Fe."
 We used a synthetic template spectrum to avoid introducing additional noise into the results., We used a synthetic template spectrum to avoid introducing additional noise into the results.
 The template was calculated using the code (Smalleyetal.2001) and adopting Το=3500K and logg=4.5., The template was calculated using the code \citep{smalleyetal01-1} and adopting $T_\mathrm{eff} = 3500~\mathrm{K}$ and $\log g = 4.5$.
" This yielded radial velocities which were less noisy but still did not demonstrate the expected variations in that they do not appear to follow the photometric period, but instead prefer P=0.323d refec|3pgrig, , bottom)."," This yielded radial velocities which were less noisy but still did not demonstrate the expected variations in that they do not appear to follow the photometric period, but instead prefer $P = 0.323~\mathrm{d}$ \\ref{ec13pg_fig}, bottom)."
 We have folded the radial velocity data on both the photometric period and the one extracted from the spectroscopic periodogram., We have folded the radial velocity data on both the photometric period and the one extracted from the spectroscopic periodogram.
" As expected, since the photometric period is barely, if at all, present in the spectroscopic data, that period yields a very poor fit refec13phrvjig, , top)."," As expected, since the photometric period is barely, if at all, present in the spectroscopic data, that period yields a very poor fit \\ref{ec13phrv_fig}, top)."
"T he""spectroscopic"" periodat firstglanceprovidesa , bottom)."," The ""spectroscopic"" period at first glance provides an acceptable fit to the data \\ref{ec13phrv_fig}, bottom)."
"However, closerinspectionrevealsthattherear« refecl"," However, closer inspection reveals that there are systematic differences between the data from the two nights, as, with one exception, the velocities from the first night all lie below the fit."
"3rvsrigwehaveplottedtheradialvelocitiesinsequenceversustime, w inedsinusoidalvariation."," In \\ref{ec13rvs_fig} we have plotted the radial velocities in sequence versus time, which makes it even more obvious that the velocities do not follow a well-defined sinusoidal variation."
Wethere f oredoubtthephysicalrelevanceo , We therefore doubt the physical relevance of this signal.
fti ," Again we point out that the longest photometric data set excludes the ""spectroscopic"" period for the light curve."
"Without more and better data, we are not able to clarify this puzzling behaviour."," Without more and better data, we are not able to clarify this puzzling behaviour."
 Perhaps it is due to a combination of the low spectral resolution and a low inclination (for the photometric variation the low inclination could be compensated for by a particularly strong reflection effect due to tees), Perhaps it is due to a combination of the low spectral resolution and a low inclination (for the photometric variation the low inclination could be compensated for by a particularly strong reflection effect due to a hot white dwarf).
 Aen1650 of this system clearly e-reso.dO Voution spectroscopy., Further investigation of this system clearly requires time-resolved high-resolution spectroscopy.
" Folding the nightly average spectra with Bessell filters we obtain V=16.26 and B-V=0.36 for April 3, and V=16.61, B—V=0.42 forApril 5."," Folding the nightly average spectra with Bessell filters we obtain $V = 16.26$ and $B\!-\!V = 0.36$ for April 3, and $V = 16.61$, $B\!-\!V = 0.42$ forApril 5."
" The difference in magnitude is very similar to that found for 112477-1738, and we attribute this and the differencein the continuum slope to the non-photometric conditions during the observations."," The difference in magnitude is very similar to that found for 12477–1738, and we attribute this and the differencein the continuum slope to the non-photometric conditions during the observations."
" Previously reported values for 113349-3237 are V= 16.34, B—V=0.36 (Kilkennyetal. 1997)."," Previously reported values for 13349–3237 are $V = 16.34$ , $B\!-\!V = 0.36$ \citep{kilkennyetal97-1}. ."
". Using the spectroscopic decomposition/fit technique introduced in , wedeterminethewhitedwar ftemperatureandmasso f —3237,Twa=35010+3415 KK, and Mya=0.46+0.11 "," Using the spectroscopic decomposition/fit technique introduced in \\ref{ec12477_sect}, , we determine the white dwarf temperature and mass of 13349--3237, $T_\mathrm{wd}=35010\pm3415$ K, and $M_\mathrm{wd}=0.46\pm0.11$ \\ref{specfit_tab} "
reduction of the field strength [rom that measured on the photosphere is needed to avoid unreasonably high Alfveen speeds. which would put too severe a limit on the time step of numerical integration.,"reduction of the field strength from that measured on the photosphere is needed to avoid unreasonably high Alfv́een speeds, which would put too severe a limit on the time step of numerical integration."
 After the smoothing. the magnetic [flux in a central area. which roughly encompasses the region of the observed flux emergence (including the rotating. positive sunspol) is zeroed oul (see Fieure lee) to be the area where the emergence of an idealized. twisted magnetic torus is driven on the lower boundary.," After the smoothing, the magnetic flux in a central area, which roughly encompasses the region of the observed flux emergence (including the rotating, positive sunspot) is zeroed out (see Figure \ref{fig1}c c) to be the area where the emergence of an idealized, twisted magnetic torus is driven on the lower boundary."
 The potential field constructed from this lower boundary normal [αν distribution in Figure lee is assumed to be the initial coronal magnetic field for our simulation. which is shown in Figure 2..," The potential field constructed from this lower boundary normal flux distribution in Figure \ref{fig1}c c is assumed to be the initial coronal magnetic field for our simulation, which is shown in Figure \ref{fig2}."
 We zero out the normal flux in (he area lor driving the [lux emergence so thal we can specily analyically the subsurface emergence structure in a field ree region without the complication of the subsurface extension of a pre-existing flux in (he same area., We zero out the normal flux in the area for driving the flux emergence so that we can specify analytically the subsurface emergence structure in a field free region without the complication of the subsurface extension of a pre-existing flux in the same area.
 The initial atmosphere in the domain is assumed to be a static polvtropic gas: where pj—8.365x10 e em.. and py=0.152 dvne cm7 are respectively the density and pressure at the lower boundary. of the coronal domain. ancl the corresponding assumed temperature al the lower boundary is 1.1. MIN.," The initial atmosphere in the domain is assumed to be a static polytropic gas: where $\rho_0 = 8.365 \times 10^{-16} $ g ${\rm cm}^{-3}$, and $p_0 = 0.152$ dyne ${\rm cm}^{-2}$ are respectively the density and pressure at the lower boundary of the coronal domain, and the corresponding assumed temperature at the lower boundary is 1.1 MK."
 The initial magnetic field in the domain is potential. and (hus does not exert anv forcing on the atmosphere which is in hyedrostatic equilibrium.," The initial magnetic field in the domain is potential, and thus does not exert any forcing on the atmosphere which is in hydrostatic equilibrium."
" Figure 3. shows the height profiles of the Allvénn speed and the sound speed along a vertical line rooted in the peak D, of the main pre-existing negative polarity spot.", Figure \ref{fig3} shows the height profiles of the Alfvénn speed and the sound speed along a vertical line rooted in the peak $B_r$ of the main pre-existing negative polarity spot.
 For (he initial state constructed. the peak Alfvénn speed is about 24 Mam/s. and the sound speed is 141 km/s at the bottom ancl gradually declines with height.," For the initial state constructed, the peak Alfvénn speed is about 24 Mm/s, and the sound speed is 141 km/s at the bottom and gradually declines with height."
 In most of the simulation domain. the Alfvénn speed is signilicantly greater than the sound speed.," In most of the simulation domain, the Alfvénn speed is significantly greater than the sound speed."
" At the lower boundary (at r= δι}. we impose (kinematically) the emergence of a twisted torus Bie by specifying a time dependent. transverse electric field E_|,2,, that corresponds to the upward advection of the torus wilh a velocity. v4: The magnetico field Τιμ,tu used for specilving$4 E_|,2,.oy is an axisvnmietrie torus defined in its own local spherical polar coordinate svstem (77. 8'. ©) whose polar axis is the svinmeltrie axis of the torus."," At the lower boundary (at $r=R_{\odot}$ ), we impose (kinematically) the emergence of a twisted torus ${\bf B}_{\rm tube}$ by specifying a time dependent transverse electric field ${\bf E}_{\perp}|_{r=R_{\odot}}$ that corresponds to the upward advection of the torus with a velocity ${\bf v}_{\rm rise}$: The magnetic field ${\bf B}_{\rm tube}$ used for specifying ${\bf E}_{\perp}|_{r=R_{\odot}}$ is an axisymmetric torus defined in its own local spherical polar coordinate system $r'$, $\theta'$, $\phi'$ ) whose polar axis is the symmetric axis of the torus."
 In the sun-centered simulation spherical coordinate svstem. the origin of," In the sun-centered simulation spherical coordinate system, the origin of"
qsPhe coexistence. of black holes. ancl starburst clusters is: known to exist in many 000galaxies. anc there are many evidences that sugeest a connection between these phenomena.,"The coexistence of black holes and starburst clusters is known to exist in many galaxies, and there are many evidences that suggest a connection between these phenomena."
 Many. studies point out that both the active nucleus ancl starbursts. might be related. to. gas. inflow. voobablv triggeredpe by an aNxis-sasvoumetry Iperturbation like bars. mergersὃν or C.tidal interactions.» (Shlosman.s Frank. Degelman; 1989. Shlosman.. ;Degelman Frank; 1990. Alaiolino et al.," Many studies point out that both the active nucleus and starbursts might be related to gas inflow, probably triggered by an axis-asymmetry perturbation like bars, mergers or tidal interactions (Shlosman, Frank Begelman 1989, Shlosman, Begelman Frank 1990, Maiolino et al."
 1997. IXnapen. Shlosman Peleticr 2000. Fathi et al.," 1997, Knapen, Shlosman Peletier 2000, Fathi et al."
 2006. ΙΟ et al.," 2006, Riffel et al."
 2008)., 2008).
 In addition. one of the most intriguing research areas in contemporary extragalactic astrophysics involves the study of the interplay between nuclear black holes. the jets. which. they can produce ane the interstellar/intergalactio. medium. (LSAT). in. which. they propagate.," In addition, one of the most intriguing research areas in contemporary extragalactic astrophysics involves the study of the interplay between nuclear black holes, the jets which they can produce and the interstellar/intergalactic medium (ISM) in which they propagate."
 EThese jets: can have a considerable: impact. on this medium., These jets can have a considerable impact on this medium.
 One aspect of jet-ISM interaction is that it can trigecr star formation., One aspect of jet-ISM interaction is that it can trigger star formation.
" Such jet-induced star formation is. considered.: a possible. mechanism. to explain.. the UVAj continuum. emission.. observed in. the host galaxies. of adistant radio sources and the ""alignment ellect between the radio emission and this continuum (ltees 1989).", Such jet-induced star formation is considered a possible mechanism to explain the UV continuum emission observed in the host galaxies of distant radio sources and the “alignment effect” between the radio emission and this continuum (Rees 1989).
 Although this ellect might play a very important role in high-z radio galaxies. detecting and studying the jet-LSAL interaction in them is very challenging.," Although this effect might play a very important role in high-z radio galaxies, detecting and studying the jet-ISM interaction in them is very challenging."
 Because of the observational, Because of the observational
"The of the CBR is eiveu by T;=To(ll:) with Ty=2,726 Ix femperature(Mather et al.",The temperature of the CBR is given by $\tr=T_0(1+z)$ with $T_0=2.726$ K (Mather et al.
 1991)., 1994).
" The evolution of the eas temperature Z, d:8 governed by the equation (see e.g. Puy et al.", The evolution of the gas temperature $\tg$ is governed by the equation (see e.g. Puy et al.
 1993: Palla et al., 1993; Palla et al.
 1995) Musa]., 1995) ].
 The first terii represents the adiabatic cooling associated with the expansion of the Universe. P being the scale factor.," The first term represents the adiabatic cooling associated with the expansion of the Universe, $R$ being the scale factor."
" The other two terms represent respectivolv the net trausfer of energy from the CBR to the eas (per unt fine and uuit voluue) via Compton scattering of CBR photons on electronsOoo amd via excitation aud de-excitation of molecular trausitious oy, where Cy; aud C; ave the collisional excitation a1 de-excitation coefficients and c; are the fractional leve populations."," The other two terms represent respectively the net transfer of energy from the CBR to the gas (per unit time and unit volume) via Compton scattering of CBR photons on electrons, and via excitation and de-excitation of molecular transitions _k where $C_{ij}$ and $C_{ji}$ are the collisional excitation and de-excitation coefficients and $x_{i}$ are the fractional level populations."
 For the molecular heating aud cooling ofthe eas we lave considered the contribution of IL. IID :ux Lill. A full discussion of the nolecular parameters is eiveu in the Appendix. where we present analytical fits iux plots of the cooling functions in the temperature ranec πα...104K.," For the molecular heating and cooling of the gas we have considered the contribution of $_2$, HD and LiH. A full discussion of the molecular parameters is given in the Appendix, where we present analytical fits and plots of the cooling functions in the temperature range $10\;{\rm K}\leq \tg \leq 10^4\;{\rm K}$."
 Given the large range of validity. th cooling functions can be used im a variety of cosinologica applications.," Given the large range of validity, the cooling functions can be used in a variety of cosmological applications."
 The energy transter funcion (IΆμωι Can become au effective heating (cooling) source for the eas if the rate of, The energy transfer function $(\Gamma-\Lambda)_{\rm mol}$ can become an effective heating (cooling) source for the gas if the rate of
in an accretion disk to explain the (win kllz-QDPOs in LMXDs.,in an accretion disk to explain the twin kHz-QPOs in LMXBs.
 Our model is able to discriminate between slow and [ast rotator as already shown in Pétri(2005a).., Our model is able to discriminate between slow and fast rotator as already shown in \cite{2005A&A...439L..27P}.
 Moreover. with help on new data from a dozen rotators. we were able to constrain (he average mass and moment of inertia of neutron stars.," Moreover, with help on new data from a dozen rotators, we were able to constrain the average mass and moment of inertia of neutron stars."
 We found for the best fit M22.0—22M. and I.220.5—L.5(10.km)?M...," We found for the best fit $M \approx 2.0-2.2 \, M_\odot$ and $I_* \approx 0.5-1.5 \,
(10\;\textrm{ km})^2 \, M_\odot$."
 Whereas the moment of inertia gives roughly the same value as (hose obtained [rom independent wavs by solving the stellar structure with several equations of state (Worleyetal.2008).. the neutron star mass appears rather large.," Whereas the moment of inertia gives roughly the same value as those obtained from independent ways by solving the stellar structure with several equations of state \citep{2008ApJ...685..390W}, the neutron star mass appears rather large."
 This effect could be an artefact of its constancy [rom one binary svstem to another., This effect could be an artefact of its constancy from one binary system to another.
 Better fits suggests to look al each svstem individually aud remove the constant mass approximation for the whole set of LAINBs. leading to a spread in the mass distiibution function for neutron stars.," Better fits suggests to look at each system individually and remove the constant mass approximation for the whole set of LMXBs, leading to a spread in the mass distribution function for neutron stars."
 Bul this requires a much more detailed separate analvsis of each binary with (heir own specilicilies (accretion rate. magnetic field strength for instance) and better observations.," But this requires a much more detailed separate analysis of each binary with their own specificities (accretion rate, magnetic field strength for instance) and better observations."
 New (ime analvzing instruments like the IITRS (hel Time Resolution Spectrometer) project on board IXO will give more insights into supra-nuclear matterand strong gravity physies (Barretetal. 2003).., New time analyzing instruments like the HTRS (High Time Resolution Spectrometer) project on board IXO will give more insights into supra-nuclear matterand strong gravity physics \citep{2008SPIE.7011E..10B}. .
straight line to some degree of accuracy.,straight line to some degree of accuracy.
 Lf we fit one to the A-band of Figure (2)). then subtract the fit. we are left with a peculiar velocity dispersion of 77 km +.," If we fit one to the $K$ -band of Figure \ref{scatter1}) ), then subtract the fit, we are left with a peculiar velocity dispersion of 77 km $^{-1}$."
" This is indistinguishable. statistically, [rom the raw dispersion: so there is no linear fit."," This is indistinguishable, statistically, from the raw dispersion; so there is no linear fit."
 In search of the most basic. qualitative signal. we then divide the plot into positive and negalive peculiar eravily halves. excluding any points whose errors would take them across the zero line.," In search of the most basic, qualitative signal, we then divide the plot into positive and negative peculiar gravity halves, excluding any points whose errors would take them across the zero line."
 For each half we form the average and estimate (he uncertaintv in the average based on the individual uncertainties in the peculiar velocities., For each half we form the average and estimate the uncertainty in the average based on the individual uncertainties in the peculiar velocities.
 For the region of positive Iv peculiar oOgravity. the averagee peculiar radial velocity is +27 km !&7: for negative.e -23 kms 1'+6.," For the region of positive $K$ peculiar gravity, the average peculiar radial velocity is +27 km $^{-1} \pm 7$ ; for negative, -23 km $^{-1} \pm 6$."
 The rms velocity dispersion for both positive and negative IX is 49 km |., The rms velocity dispersion for both positive and negative K is 49 km $^{-1}$.
 But again. the uncertainties in peculiar velocity are correlated. so the stated errors could be very misleading: and the dispersions are much larger than the averages in magnitude.," But again, the uncertainties in peculiar velocity are correlated, so the stated errors could be very misleading; and the dispersions are much larger than the averages in magnitude."
 There could be a general. average correlation of peculiar velocity. with svnthetüe gravity. but at best il explains a minoritv of the actual motion. aud it is not certain il exists.," There could be a general, average correlation of peculiar velocity with synthetic gravity, but at best it explains a minority of the actual motion, and it is not certain it exists."
 It is time {ο examine the model itself in terms of its dvnamical implications., It is time to examine the model itself in terms of its dynamical implications.
 Up to this time we have used the average. kinematicallv-derived background model as a basis for peculiar velocities. and have found no correlation between velocities (apples) ancl Iuninositv-derived peculiar gravity. (oranges).," Up to this time we have used the average, kinematically-derived background model as a basis for peculiar velocities, and have found no correlation between velocities (apples) and luminosity-derived peculiar gravity (oranges)."
 Suppose we (rv a comparison which adjusts (he model to mininize the difference between the peculiar velocity ancl the A-band svuthetic gravilv: We are [aced with the fact (hat the normalization of the svuthetic gravity. g;. is unknown.," Suppose we try a comparison which adjusts the model to minimize the difference between the peculiar velocity and the $K$ -band synthetic gravity: We are faced with the fact that the normalization of the synthetic gravity, $g_i$, is unknown."
 If it is too small. we are essentially reproducing the kinematic model: if it is too big. we are filling a model to the Iuminositv field and ignoring motionsa caleulation which might be of some interest. but not to us now.," If it is too small, we are essentially reproducing the kinematic model; if it is too big, we are fitting a model to the luminosity field and ignoring motions—a calculation which might be of some interest, but not to us now."
 In practice a series of models was calculated with different normalizations. starting with one which matched the rms value of the peculiar velocity in (he isotropic kinematic model.," In practice a series of models was calculated with different normalizations, starting with one which matched the rms value of the peculiar velocity in the isotropic kinematic model."
 This turned out in fact to give the best correlation. shown in theleft-hand panel of Figure," This turned out in fact to give the best correlation, shown in theleft-hand panel of Figure"
"where (0,..0,) is its angular position ou the sky. aud A, and dy are plauar cocfiicicuts for the leloccutric velocity distribution of the NDP.","where $(\theta_x,\theta_y)$ is its angular position on the sky, and $A_{x}$ and $A_{y}$ are planar coefficients for the heliocentric velocity distribution of the KDP."
 This equation is simular to equation (6)) but we have replaced the Ac in the NDP terius bv ον ie. we fit to the heloceutric rather than the residual velocities.," This equation is similar to equation \ref{twopopp}) ) but we have replaced the $\Delta v$ in the KDP terms by $v$, i.e., we fit to the heliocentric rather than the residual velocities."
" The origen of our x-x coordinate svsteni jj ata52]"".8—GOPLT. with X increasing to the east aud Y to the north."," The origen of our x-y coordinate system is at $\alpha=5^h21^m, \delta= -69^\circ17^\prime$, with X increasing to the east and Y to the north."
" ήΠαν we set A,=0. so that there are same munber of deerees of freedom as in the three-Craussian fit to the residuals."," 	Initially, we set $A_{x}=A_{y}=0$, so that there are same number of degrees of freedom as in the three-Gaussian fit to the residuals."
 We find no solution here that has a lower P than the two-Caussian solution. implying that there is uo evidence for the existence of a third population having a conunon hehocentric velocity outside the LAIC disk.," We find no solution here that has a lower $\chi^2$ than the two-Gaussian solution, implying that there is no evidence for the existence of a third population having a common heliocentric velocity outside the LMC disk."
" We therefore repeat the search. but allow iL, and <4, to vary as free parauicters."," 	We therefore repeat the search, but allow $A_{x}$ and $A_{y}$ to vary as free parameters."
 We find that the likelihood is then maximized at very low values of the velocity dispersion okpp<Ikus!.," We find that the likelihood is then maximized at very low values of the velocity dispersion $\sigma_{\rm KDP}\la 1\,\kms$."
 We reject these soultious as uulivsical. and note that our fitting routines iav have been falsely attracted to them as results of inevitable Poisson noise.," We reject these soultions as unphysical, and note that our fitting routines may have been falsely attracted to them as results of inevitable Poisson noise."
" We then find a solution with 39 stars in the KDP with eypbp=lW6.tkmsLf. A,=δανtdeg |, A,=L9lamnstdee to oppp=Shans aud Gxpp=0.673."," 	We then find a solution with 39 stars in the KDP with $\bar v_{KDP}
= 16.4\,\kms$, $A_x = 2.6\,\kms\,\rm deg^{-1}$ , $A_y = 4.9\,\kms\,\rm
deg^{-1}$ , $\sigma_{\rm KDP}=5\,\kms$ , and $\zeta_{\rm KDP}=0.673$."
 Relative to the two-Caussian solution. this KDP solution has A4?=16 for 6 additional parameters.," Relative to the two-Gaussian solution, this KDP solution has $\Delta\chi^2=16$ for 6 additional parameters."
 Figure 6. shows the residuals of the LAIC stars with respect to the NDP., Figure \ref{kdpfig} shows the residuals of the LMC stars with respect to the KDP.
 The KDP is shown as the strong peak of points around residual 0., The KDP is shown as the strong peak of points around residual 0.
 Other small peaks are due to the clumped distribution of our stars in angle. aud are not sienificaut.," Other small peaks are due to the clumped distribution of our stars in angle, and are not significant."
 There are not enough stars in the NDP to siguificautlv determine if the NDP covers the cutive face of the LMC or has a patchy distribution., There are not enough stars in the KDP to significantly determine if the KDP covers the entire face of the LMC or has a patchy distribution.
 While the probability that amy randomly chosen plane will come within ~5laus| of a significant fraction of our suuple stars is sinall (aud well represcuted by the \? test). there are a laree nuuber of iudependeut planes that can be compared to the data.," 	While the probability that any randomly chosen plane will come within $\sim 5\,\kms$ of a significant fraction of our sample stars is small (and well represented by the $\chi^2$ test), there are a large number of independent planes that can be compared to the data."
 To obtain a more accurate assesslucut of the statistical significance of this detection. we perform a set of Monte Carlo simulations.," To obtain a more accurate assessment of the statistical significance of this detection, we perform a set of Monte Carlo simulations."
 Iu cach siuulation. we draw velocities randomly from the Caussian distribution of disk residuals found im 33.," In each simulation, we draw velocities randomly from the two-Gaussian distribution of disk residuals found in 3."
 We then search for a NDP in the resulting helioceuntric velocities iu the same way we did for the actual data in 11.1 and 11.2., We then search for a KDP in the resulting heliocentric velocities in the same way we did for the actual data in 4.1 and 4.2.
 In order to make the simulations tractable. we ignore metallicity information.," In order to make the simulations tractable, we ignore metallicity information."
" This simplification is justified by the fact that the metallicity of the NDP measured in 11.2 is not siguiicantlv different from the ""young disk component.", This simplification is justified by the fact that the metallicity of the KDP measured in 4.2 is not significantly different from the “young disk” component.
 If metallicity is ignored then the external-plane solution⋅ shows an improvement⋅ of. Αντ5=111 forE 5 additional paraiucters. which is formally significant at the level.," If metallicity is ignored then the external-plane solution shows an improvement of $\Delta\chi^2=14$ for 5 additional parameters, which is formally significant at the level."
 However. we find that out of LOT simulations. there is AQ?>Ll in 26 cases.," However, we find that out of 407 simulations, there is $\Delta\chi^2\geq 14$ in 26 cases."
 Hence. our detection is significant oulv at the level. roughly equivalent to σσ.," Hence, our detection is significant only at the level, roughly equivalent to $2\,\sigma$."
 Given the intriguing signal we see in the C star velocities. but also the mareinal level of significance. it is worth exploring other possible signs of the NDP.," Given the intriguing signal we see in the C star velocities, but also the marginal level of significance, it is worth exploring other possible signs of the KDP."
 One such tracer is the 21c01à gas cussion. mapped. ce. bv Lids BRohlfs (1992)). and ἵνα (1998)).," One such tracer is the 21cm gas emission, mapped, e.g., by Luks Rohlfs \cite{lh}) ), and Kim \cite{kim}) )."
" Luks Rohlfs note that a lower velocity component (""L-conmponeut) coutains about of the IIT gas in the LMC. is separated from the main velocity component bv 30 kin/s. Although Wim (1998)) do not specifically conunent on such a compoucut iu their paper based on lugher spatial resolution WT inagiug. a similar sigual secs evident iu their position-velocity maps (c.¢.. Ta and Th in their paper) at RA 05:37 - 05:17 and DEC -30 to -120 arcmin."," Luks Rohlfs note that a lower velocity component (``L-component'') contains about of the HI gas in the LMC, is separated from the main velocity component by $\sim$ 30 km/s. Although Kim \cite{kim}) ) do not specifically comment on such a component in their paper based on higher spatial resolution HI imaging, a similar signal seems evident in their position-velocity maps (e.g., 7a and 7b in their paper) at RA 05:37 - 05:47 and DEC -30 to -120 arcmin."
 The standard interpretation of this substructure iu gas is that it is due to hydrodynamic effects on gas within the LAIC disk., The standard interpretation of this substructure in gas is that it is due to hydrodynamic effects on gas within the LMC disk.
 However. the correlation of the gas velocity “L-component™” with the stellar IKDP sugecsts that the eas may be outside the LAIC disk.," However, the correlation of the gas velocity “L-component” with the stellar KDP suggests that the gas may be outside the LMC disk."
 An intriguing but somewhat more ambiguous signature niv be evident in the CID star velocities of Cowley ILutwick (1991))., An intriguing but somewhat more ambiguous signature may be evident in the CH star velocities of Cowley Hartwick \cite{ch}) ).
 Velocities for a sample of —s0 CTI stars show a low velocity asvuuuetric tail. consisteut with a component at ~20 kin/sec lower svstematic velocity.," Velocities for a sample of $\sim$ 80 CH stars show a low velocity asymmetric tail, consistent with a component at $\sim$ 20 km/sec lower systematic velocity."
 Cowley Tartwick (1991) even sugeest that one explanation of this population is that it is a result of an earlier violent tidal encouuter between the LAIC-SAIC system and the Milev Wary., Cowley Hartwick (1991) even suggest that one explanation of this population is that it is a result of an earlier violent tidal encounter between the LMC-SMC system and the Milky Way.
 The small sample statistics aud asviuuetrice spatial distribution of these stars make a more detailed exploration difficult., The small sample statistics and asymmetric spatial distribution of these stars make a more detailed exploration difficult.
 Wo mav have detected. a kinematically distinct population of carbon stars in the direction of the LAIC., We may have detected a kinematically distinct population of carbon stars in the direction of the LMC.
 If real. this population could be either a structure witlin the LMC disk or tidal debris that is well separated from the disk aud lence either in front of or behind the LMC.," If real, this population could be either a structure within the LMC disk or tidal debris that is well separated from the disk and hence either in front of or behind the LMC."
 If it is well separated from the LMC. then it would eive vise to microlensine: either it would be iu front of the LMC and so would act as lenses. or it would be behind the LAIC and would act as sources.," If it is well separated from the LMC, then it would give rise to microlensing: either it would be in front of the LMC and so would act as lenses, or it would be behind the LMC and would act as sources."
 The iicrolensing optical depth due to a thin sheet of stellar matter with deusitv X4 aud the LAIC with deusitv Xo separated by a distance D which is small compared to the distance from the Sun to the LMC is: The distance between the two sheets. D. caunot be determined from velocity data alone.," The microlensing optical depth due to a thin sheet of stellar matter with density $\Sigma_1$ and the LMC with density $\Sigma_2$ separated by a distance $D$ which is small compared to the distance from the Sun to the LMC is: The distance between the two sheets, $D$, cannot be determined from velocity data alone."
 However. since the two sheets umust have similar velocities. the tidal tail cannot be a random interloper iu the halo. but must be somehow related to the LAC.," However, since the two sheets must have similar velocities, the tidal tail cannot be a random interloper in the halo, but must be somehow related to the LMC."
 Lacking further information. we amake thesomewhat ad hoc assumption that the material in the tidal tail has been moving awayfrou the LAIC at a constant velocity of 30;ans| suce close tidal," Lacking further information, we make thesomewhat ad hoc assumption that the material in the tidal tail has been moving awayfrom the LMC at a constant velocity of $30\,
\kms$ since close tidal"
and Nay flux can lead to enhanced ICN starting by the ionization of No.,and X-ray flux can lead to enhanced HCN starting by the ionization of $_2$.
" The columm deusities for IICN aud CS from the Stauberetal. model of protostar AFGL 2591 agree with our observed. abundanuces while their cobi densities forCIT»...CIT, aud aare lower than our observed values."," The column densities for HCN and CS from the \citeauthor{stauber05}
 model of protostar AFGL 2591 agree with our observed abundances while their column densities for, and are lower than our observed values."
 Tot core aud disk chenüstrv iuodels predict the Chhancement of molecules such asColle...CIT... aud NII5.," Hot core and disk chemistry models predict the enhancement of molecules such as, and ."
. However. the observed abuudauce of iis higher than what models predict frou wii gas-phase chemustry (see Table 53).," However, the observed abundance of is higher than what models predict from warm gas-phase chemistry (see Table \ref{tab:mods}) )."
" This iudicates that lis probably frozen ou dust erains aud subliamates along with molecules likeCIT,.", This indicates that is probably frozen on dust grains and sublimates along with molecules like.
. Boudinetal.(1998). studied the solid features ofCol... especially when mixed with oor CO.," \citet{boudin98} studied the solid features of, especially when mixed with or CO."
 They found that the features broaden substantially when mixed with aand to a lesser extent when mixed with CO., They found that the features broaden substantially when mixed with and to a lesser extent when mixed with CO.
 Thev compare the laboratory data to ddata for the colder neighbor. IRS 9. which resulted iu an upper limit of 8&LOL ffor the column deusitv of solidCo.," They compare the laboratory data to data for the colder neighbor, IRS 9, which resulted in an upper limit of $8 \times 10^{17}$ for the column density of solid."
.. This upper Init is consistent with the observed gas-phase column deusity seen toward1., This upper limit is consistent with the observed gas-phase column density seen toward.
. So. it is possible that solid lis sublianating from erain mautles as protostars heat the enviroment.," So, it is possible that solid is sublimating from grain mantles as protostars heat the environment."
" The observed irafio agrees with the branching ratio for the destruction ofCIT,.", The observed ratio agrees with the branching ratio for the destruction of.
. Based on these results «πλοία also be very abunudaut., Based on these results should also be very abundant.
 We now attempt fo construct a physical aud econietrical iiodel ofix., We now attempt to construct a physical and geometrical model of.
 Within the possible scenarios prescuted from various radio and infrared observations (e.g...Minierrausetal. 2006).. we propose a scenario dn which the molecular absorption preseuted here comes from a circustellar disk.," Within the possible scenarios presented from various radio and infrared observations \citep[e.g.,][]{minier01, lugo04,
debuizer05, kraus06}, we propose a scenario in which the molecular absorption presented here comes from a circumstellar disk."
 We will cxamine other possibilities fist., We will examine other possibilities first.
 Figure 8. depicts the possible scenarios allowed by the available observations., Figure \ref{fig:cartoon} depicts the possible scenarios allowed by the available observations.
 We can quickly rule out a simple model in which the absorbing molecular gas is in the foreground molecular cloud auc not closelv associated with1., We can quickly rule out a simple model in which the absorbing molecular gas is in the foreground molecular cloud and not closely associated with.
" Our observations require the gasto be much hotter (T~ 300 Is) aud deuser Gn,10* tto maintain rotational LTE of ICN out to 7= 21) than is found away from Iuninous sources in molecular clouds.", Our observations require the gas to be much hotter $T \sim$ 300 K) and denser $n_{\rm H_2} \sim 10^7$ to maintain rotational LTE of HCN out to $J = 21$ ) than is found away from luminous sources in molecular clouds.
 According to vanderTaketal.(2000)... teniperatures do uot reach the observed values uutil vou ect to within 100 AU of the ceutral star.," According to \citet{vdt00}, temperatures do not reach the observed values until you get to within 400 AU of the central star."
" Iu addition. the radiation field ust have a brieltuess teniperature ~300 I& to populate the vs, level sufficiently to account for the observed aabsorptiou."," In addition, the radiation field must have a brightness temperature $\sim$ 300 K to populate the $\nu_5$ level sufficiently to account for the observed absorption."
 Undoubtedly. the absorbing molecular gas is in close proximity to the ccontinuuii source1.," Undoubtedly, the absorbing molecular gas is in close proximity to the continuum source."
. Tt is not quite so easy to rule out a model in which the absorbing molecules are in boundary region between the euvelope around the hhypercompact II II region aud the outflow., It is not quite so easy to rule out a model in which the absorbing molecules are in boundary region between the envelope around the hypercompact H II region and the outflow.
 This eas could be compressed by the ionized wind. aud if it is as close as0.1. ov 280 AU. from a 10LE. source it would have a temperature near 300 IN. The temperature structure cletermined by vanderTaketal.(2000) iudicates that indecd temperatures range from 200 £00 I at radii of 220380 AU.," This gas could be compressed by the ionized wind, and if it is as close as, or 280 AU, from a $10^{5}~L_{\odot}$ source it would have a temperature near 300 K. The temperature structure determined by \citet{vdt00} indicates that indeed temperatures range from 200 – 400 K at radii of 220–380 AU."
 However. interaction with the ionized wind. which has a velocity ~100 |. would be expected to accelerate the gas. causing broad. bluc-shifted absorption (seevanderTaketal.2000).," However, interaction with the ionized wind, which has a velocity $\sim$ 100, would be expected to accelerate the gas, causing broad, blue-shifted absorption \citep[see][]{vdt00}."
. Iu contrast. the observed lines have widths of <8 aand have ceutroids within a few oof those seen in surrounding molecular eas.," In contrast, the observed lines have widths of $<8$ and have centroids within a few of those seen in surrounding molecular gas."
 Another argunent agaimst the preseuce of the observed eas in au envelope around lis the fact that uct absorption is seen., Another argument against the presence of the observed gas in an envelope around is the fact that net absorption is seen.
 Since the vibrational temperature is comparable to the brightuess temperature of the coutimmiun radiation. cussion lines would be seen if the molecular gas had a larger exteut than the continua source.," Since the vibrational temperature is comparable to the brightness temperature of the continuum radiation, emission lines would be seen if the molecular gas had a larger extent than the continuum source."
 If the lines arose in 1iolecular shell surrounding the Lypercompact IT II region. this probably would be the case.," If the lines arose in molecular shell surrounding the hypercompact H II region, this probably would be the case."
 Ou the other haud. Campbell(1981). proposes that the centimeter contiumuunau from IRS 1 results from partially ionized material in an outflow.," On the other hand, \citet{campbell84} proposes that the centimeter continuum from IRS 1 results from partially ionized material in an outflow."
 The centimeter continu cluission has been spatially resolved iuto knots by Caeetal. (1995).. who sugecst that the emission comes from photo-evaporation of kuots of neutral molecular material (see Fie. 8)).," The centimeter continuum emission has been spatially resolved into knots by \citet{gaume95}, , who suggest that the emission comes from photo-evaporation of knots of neutral molecular material (see Fig. \ref{fig:cartoon}) )."
 The nuüniasers seen by Minieretal.(2000) also ποσα to trace the knotty structure (iu addition to the disk described im 813) probed by the cin coutimmiun cussion., The masers seen by \citet{minier00} also seem to trace the knotty structure (in addition to the disk described in \ref{sec:intro}) ) probed by the cm continuum emission.
 The knots may be material stripped fromthe disk., The knots may be material stripped fromthe disk.
 The stellar, The stellar
 high-2 , ${\it z}$ 
Here. Ay. Ao. aud Ax are expressions involving(e fy aud g4: From the 0-componeut. equation (16)). we obtain a second relation between fy aud gοἱ where we have delined0.. and where the three terms on the rightliand side are again combinations of fj and g4: We have already found. { and g4 in the subsonie case of interest.,"Here, ${\cal A}_1$ , ${\cal A}_2$ , and ${\cal A}_3$ are expressions involving $f_1$ and $g_{-1}$ : From the $\theta$ -component, equation \ref{eqn:eulert}) ), we obtain a second relation between $f_0$ and $g_{-2}$: where we have defined, and where the three terms on the righthand side are again combinations of $f_1$ and $g_{-1}$: We have already found $f_1$ and $g_{-1}$ in the subsonic case of interest."
 After substituting these expressions. equations (25)) auc (26)). iuto the righthaud sides of equatious (27)) aud (31)). the coupled equatious for fij aud g2 become: and These last two relatious constitute oursecond-order equatious.," After substituting these expressions, equations \ref{eqn:gm1}) ) and \ref{eqn:f1})), into the righthand sides of equations \ref{eqn:secondr}) ) and \ref{eqn:secondt}) ), the coupled equations for $f_0$ and $g_{-2}$ become: and These last two relations constitute our equations."
 For auy value of 3. we may integrate them uumerically [rom the upstream axis.s.. to the downstream axis atQ.," For any value of $\beta$, we may integrate them numerically from the upstream axis, to the downstream axis at."
. Three initial coucitious are required. of which we lave already identified two:.," Three initial conditions are required, of which we have already identified two:."
. As a third initial condition. we use g((8). whose value at this point is arbitrary.," As a third initial condition, we use $g_{-2} (\pi)$, whose value at this point is arbitrary."
 For each chosen value of gy»(x). we may find f/o(0) aud gο(0).," For each chosen value of $g_{-2} (\pi)$, we may find $f_0 (\theta)$ and $g_{-2} (\theta)$."
 We thus have a one-parameter family of outer flow solutionW., We thus have a one-parameter family of outer flow solutions.
 h[un he upper panel ofFigure 3. we display. for the representative value0.5.. tliree solutions of g3(0).," In the upper panel ofFigure \ref{fig:2ndfandg} we display, for the representative value, three solutions of $g_{-2} (\theta)$."
 We obtained each solution by asstuning dillerent values of gà(x)., We obtained each solution by assuming different values of $g_{-2} (\pi)$.
 Notice that g’5 vanishes ou the upstream aud downstream axes. implviug again that the density profile is [lat in bothregions.," Notice that $g_{-2}^\prime$ vanishes on the upstream and downstream axes, implying again that the density profile is flat in bothregions."
Notice also that all curves attaiu the same value at πο.,Notice also that all curves attain the same value at .
 That is. g3(2/2) depends only ou 2. auc not ou theprescribed initial condition g 2(7).," That is, $g_{-2} (\pi/2)$ depends only on $\beta$ , and not on theprescribed initial condition $g_{-2} (\pi)$ ."
orders of magnitude smaller than that found. for optically-selected. galaxies (e.g. Huterer. Knox Nichol 2001): this is readily explained by the wide redshift range of the NVSS sources. which vastly dilutes the clustering signal through the superposition of unrelated redshift slices.,"orders of magnitude smaller than that found for optically-selected galaxies (e.g. Huterer, Knox Nichol 2001); this is readily explained by the wide redshift range of the NVSS sources, which vastly dilutes the clustering signal through the superposition of unrelated redshift slices."
" The NVSS signal remains 5 orders of magnitude greater than the CAIB €, spectrum over the same multipole range. reflecting the erowth of structure since z=1100."," The NVSS signal remains $\sim 5$ orders of magnitude greater than the CMB $C_\ell$ spectrum over the same multipole range, reflecting the growth of structure since $z = 1100$."
" Given the incomplete. sky and. finite resolution. the measured C, values are not independent."," Given the incomplete sky and finite resolution, the measured $C_\ell$ values are not independent."
 However. it was argued in Section 3.2. that the correlations between neighbouring power spectrum measurements are small.," However, it was argued in Section \ref{secestharm} that the correlations between neighbouring power spectrum measurements are small."
 Phis fact was confirmed by the maximum likelihood analysis., This fact was confirmed by the maximum likelihood analysis.
 Phe degree of correlation between neighbouring bins is given by the immecdiately oll-diagonal elements of the inverse Fisher matrix., The degree of correlation between neighbouring bins is given by the immediately off-diagonal elements of the inverse Fisher matrix.
 This data is generated by the ALADCAD software. and inspection revealed that the size of the immediately olf-diagonal matrix elements was 20 times smaller than that of the diagonal elements.," This data is generated by the MADCAP software, and inspection revealed that the size of the immediately off-diagonal matrix elements was $\sim 20$ times smaller than that of the diagonal elements."
 Figure 7 compares the angular power spectra measured at [lux-density thresholds 5 mJy. LO my and 20 my using spherical harmonic analysis.," Figure \ref{figclflux} compares the angular power spectra measured at flux-density thresholds 5 mJy, 10 mJy and 20 mJy using spherical harmonic analysis."
 The 5 mJv data may be allected bv svstematic surface density gradients., The 5 mJy data may be affected by systematic surface density gradients.
 The results are consistent with an unchanging underlying power spectrum., The results are consistent with an unchanging underlying power spectrum.
 This is not surprising: the redshift. distribution of radio sources does not vary significantly between 5 my. ancl 20 mJy and the angular correlation function has been found not to depend on lux density in this range (Blake Wall 2002a)," This is not surprising; the redshift distribution of radio sources does not vary significantly between 5 mJy and 20 mJy, and the angular correlation function has been found not to depend on flux density in this range (Blake Wall 2002a)."
 An interesting SIprobe of the 8galaxy 1pattern is the distribution ol values. of Ain|? (see Hauser Pechles 1973)., An interesting probe of the galaxy pattern is the distribution of values of $|A_{\ell m}|^2$ (see Hauser Peebles 1973).
 These «quantities are measured as part of our spherical harmonic analvsis (Section 3.2))., These quantities are measured as part of our spherical harmonic analysis (Section \ref{secestharm}) ).
 For a random distribution. with surface density. ay over a full sky. the central limit theorem ensures. that the real anc imaginary parts of kyneoν»»nn(7) are clrawn independentlv from. Gaussian distributionsZi such that the normalization satisfies alnn ay.," For a random distribution with surface density $\sigma_0$ over a full sky, the central limit theorem ensures that the real and imaginary parts of $A_{\ell m} = \sum_i Y_{\ell
m}^*(i)$ are drawn independently from Gaussian distributions such that the normalization satisfies $|A_{\ell m}|^2 = \sigma_0$ ."
" lt is then easy to show that vw=al,|?"" has an exponential probability distribution for nmz0: For à partial. sky. Ans]2 is replaced. by tudes|Ji, (equation 5))."," It is then easy to show that $x = |A_{\ell m}|^2$ has an exponential probability distribution for $m \ne 0$: For a partial sky, $|A_{\ell m}|^2$ is replaced by $|A_{\ell m} -
\sigma_0 I_{\ell m}|^2/J_{\ell m}$ (equation \ref{eqclpeeb}) )."
 ligure 8. plots the distribution of observed: values of Adesmodosως , Figure \ref{figalmdist} plots the distribution of observed values of $|A_{\ell m} - \sigma_0 I_{\ell m}|^2/J_{\ell m}$.
We restrict this plot to the multipole range 51«/(100: for this range of f. Ligure 5 demonstrates that C;zc0 and thus the survey is. well-described by a random clistribution with additional multiple components.," We restrict this plot to the multipole range $51 < \ell < 100$: for this range of $\ell$, Figure \ref{figcl} demonstrates that $C_\ell \approx 0$ and thus the survey is well-described by a random distribution with additional multiple components."
 For cach £ we included the range 1xmxf (negative values of m are not independent)., For each $\ell$ we included the range $1 \leq m \leq \ell$ (negative values of $m$ are not independent).
 Overplotted on Figure ὃ as the solid line is the prediction. of equation 14.., Overplotted on Figure \ref{figalmdist} as the solid line is the prediction of equation \ref{eqalmdist}.
 Multiple components cause the slope of the observed exponential distribution to be shallower than this prediction., Multiple components cause the slope of the observed exponential distribution to be shallower than this prediction.
" Section 3.3 shows that the value of <chinσυτω""hd> is increased from ay to (1|Pejou. where e is the fraction of galaxies split into double sources."," Section \ref{secmult} shows that the value of $<|A_{\ell m} - \sigma_0 I_{\ell m}|^2/J_{\ell m}>$ is increased from $\sigma_0$ to $(1+2e)\sigma_0$, where $e$ is the fraction of galaxies split into double sources."
 Thus equation 14 must be amended such that Gr)xexpμα|2e)o0]., Thus equation \ref{eqalmdist} must be amended such that $P(x) \propto \exp{[-x/(1+2e)\sigma_0]}$.
 Assuming that e=0.07. this corrected. prediction is plotted on Figure S as the dashed line and provides a very good fit to the observed distribution.," Assuming that $e = 0.07$, this corrected prediction is plotted on Figure \ref{figalmdist} as the dashed line and provides a very good fit to the observed distribution."
 This is an independent demonstration that approximately 7 per cent of NVSS galaxies are split into multiple-component sources., This is an independent demonstration that approximately 7 per cent of NVSS galaxies are split into multiple-component sources.
 Figure S also underlines the fact that the imprint of clustering on the projected radio skv is very faint., Figure \ref{figalmdist} also underlines the fact that the imprint of clustering on the projected radio sky is very faint.
 The angular correlation function. w(8) has been measured for the NVSS by Blake Wall (20022) anc Overzier ct al. (, The angular correlation function $w(\theta)$ has been measured for the NVSS by Blake Wall (2002a) and Overzier et al. (
2003).,2003).
 It is well-described bx a power-law (0)z(1 for angles up to a few degrees.," It is well-described by a power-law $w(\theta) \approx (1 \times 10^{-3}) \,
\theta^{-0.8}$ for angles up to a few degrees."
 Equation 3. allows us to derive the equivalent C5 spectrum if we assume that this power-law extends to all angular scales.," Equation \ref{eqwtocl}
 allows us to derive the equivalent $C_\ell$ spectrum if we assume that this power-law extends to all angular scales."
 In Figure 9. we overplot the resulting prediction on the measurements and find an excellent. fit.," In Figure \ref{figwth}
 we overplot the resulting prediction on the measurements and find an excellent fit."
 This is initially surprising: the angular correlation, This is initially surprising: the angular correlation
dynnamics.,namics.
 11904., 1904.
 oobtained., obtained.
 cllusters., lusters.
 staar dynamics, ar dynamics
and nuclear reactions can substantially affect the productivity of p-nuclei.,and nuclear reactions can substantially affect the productivity of p-nuclei.
 This iuakes it difficult to determine the role of the mp-process as the source of the solar p-uuclei., This makes it difficult to determine the role of the $\nu$ p-process as the source of the solar p-nuclei.
" Weeping such uucertainties iu ή, we discuss a possible coutribution of the rp-process to the solar p-abundauces based. on our result bv comparing with other possible sources."," Keeping such uncertainties in mind, we discuss a possible contribution of the $\nu$ p-process to the solar p-abundances based on our result by comparing with other possible sources."
 Table 1 lists the currently proposed. astrophysical orieius for each p-nuclde κ columu) with its solar abundance aud fraction relative to its elemental abuudanee (2ndand3rdcolumms.Lodders2003).," Table 4 lists the currently proposed astrophysical origins for each p-nuclide (1st column) with its solar abundance and fraction relative to its elemental abundance \citep[2nd and 3rd
columns,][]{Lodd2003}."
" All these sources are associated with core-collapse supernovac,", All these sources are associated with core-collapse supernovae.
Photo-dissociation of pre-existing neutrou-rich abuudances in the oxvecn-neon laver of core-collapse supernovae (or m their pre-collapse phases). ic. tle 5- process (Woosley&Toward1978:Prautzosetal.1990:al.2008) is currently regarded as the most successful scenario.,"Photo-dissociation of pre-existing neutron-rich abundances in the oxygen-neon layer of core-collapse supernovae (or in their pre-collapse phases), i.e., the $\gamma$ -process \citep{Woos1978, Pran1990, Raye1995, Raus2002,
 Haya2008} is currently regarded as the most successful scenario."
" Iu the Hth column of Table 1. the p-nuclei whose oriems cau be explained by the 5-process in Ravetctal.(1095) are specified by ""wes."," In the 4th column of Table 4, the p-nuclei whose origins can be explained by the $\gamma$ -process in \citet{Raye1995} are specified by “yes”."
 The bracketed. oues are hose underproduced in a more recent work by etal. (20023., The bracketed ones are those underproduced in a more recent work by \citet{Raus2002}.
. The origins of up to 21 out of 35 p-isotopes can be explained by the 5-process., The origins of up to 24 out of 35 p-isotopes can be explained by the $\gamma$ -process.
" However. the light isotopes ολο, PORPSRa, ρα, 106108(80, H5 Dy. ay ""Suy which account for a laree fraction iu the solar abundances. and some heavy p-isotopes (La. iix 7260) need other sources (specified by ""no in Table I)."," However, the light p-isotopes $^{92, 94}$ Mo, $^{96, 98}$ Ru, $^{102}$ Pd, $^{106, 108}$ Cd, $^{113}$ In, and $^{115}$ Sn), which account for a large fraction in the solar p-abundances, and some heavy p-isotopes $^{138}$ La and $^{152}$ Gd) need other sources (specified by “no” in Table 4)."
 The r-process (SthcolumninTableLWoosleyetal.1990) in core-collapse supernovae is suggestec o account for the production of a couple of j)eavyo pdsotopes La and UUTa (the former is nuderproduced ia the +-process).," The $\nu$ -process \citep[5th column in Table~4,][]{Woos1990} in core-collapse supernovae is suggested to account for the production of a couple of heavy p-isotopes $^{138}$ La and $^{180}$ Ta (the former is underproduced in the $\gamma$ -process)."
" The a-rvich aac slightly ueutronaich (35z0.17—0.19: slightly more xoton-rich than the stabilitv values) neutrino- outflows were also suggested as the production site ofsome light p-isotopes including ""Mo (butnotprocess.91", The $\alpha$ -rich and slightly neutron-rich $Y_\mathrm{e} \approx 0.47-0.49$; slightly more proton-rich than the $\beta$ -stability values) neutrino-driven outflows were also suggested as the production site ofsome light p-isotopes including $^{92}$ Mo \citep[but not $^{94}$ .
 The protou-richuess relative to the .+stability liue in the fragmented QSE clusters (Ποιαetal.1996:Moveretal.105) at Ty—L23 leads to the formation of these p-uuclei with /Nxi50.," The proton-richness relative to the $\beta$ -stability line in the fragmented QSE clusters \citep{Hoff1996,
 Meye1998b} at $T_9 \sim 4-3$ leads to the formation of these p-nuclei with $N \le 50$."
" Such OSE clusters on the proton-rich side of the o-stabilitv line will be denoted as ""p-QSE"" hereafter.", Such QSE clusters on the proton-rich side of the $\beta$ -stability line will be denoted as “p-QSE” hereafter.
" A recent study of nucleosvuthesis in the electrou-capture supernovac of a 9AL. star shows that the lightest p-nuclei *!Se; Kr. του, and ""?Mo can be produced in p-QSE enough to account for their solar amounts (6thcolumninTable[.Wanajooetal. 2000)."," A recent study of nucleosynthesis in the electron-capture supernovae of a $9\, M_\odot$ star shows that the lightest p-nuclei $^{74}$ Se, $^{78}$ Kr, $^{84}$ Sr, and $^{92}$ Mo can be produced in p-QSE enough to account for their solar amounts \citep[6th column in
Table~4,][]{Wana2009}."
" Tlowever. these additional sources still cannot fll the eap for some light p-isotopes such as ?!Mo, 99? Πτι, IU2pq. 106.108 C] οι, 15S). and for a heavy. p-isotope L9."," However, these additional sources still cannot fill the gap for some light p-isotopes such as $^{94}$ Mo, $^{96,
 98}$ Ru, $^{102}$ Pd, $^{106, 108}$ Cd, $^{113}$ In, $^{115}$ Sn, and for a heavy p-isotope $^{152}$ Gd."
 Qur result in this study is based on a seii-analvtic model of ueutrino-driven winds. while the results for the *-process. the r-process. aud the p-QSE listed in Table tare all based on realistic hvdrodyvuauic studies.," Our result in this study is based on a semi-analytic model of neutrino-driven winds, while the results for the $\gamma$ -process, the $\nu$ -process, and the p-QSE listed in Table 4 are all based on realistic hydrodynamic studies."
 Nevertheless. we attempt to present a list of the p-isotopes Whose origin cau be attributed to the vp-process. as follows.," Nevertheless, we attempt to present a list of the p-isotopes whose origin can be attributed to the $\nu
p$ -process, as follows."
 The requisite overproduction factor for a even nuclidececent. which explains its solar origin. is inferred to be 2LO (e...Woosleyctal.199," The requisite overproduction factor for a given nuclide, which explains its solar origin, is inferred to be $> 10$ \citep[e.g.,][]{Woos1994}."
 Asstuning the masses of the total ejecta aud of the 1)...neutrino-driven ejecta to be —104. aud ~ (e.g...Wanajo2006). the overproduction factor per supernova event is diluted by about orders of magnitude compared to our result.," Assuming the masses of the total ejecta and of the neutrino-driven ejecta to be $\sim 10\, M_\odot$ and $\sim 10^{-3}\, M_\odot$ \citep[e.g.,][]{Wana2006}, the overproduction factor per supernova event is diluted by about 4 orders of magnitude compared to our result."
 We thus apply the condition f>LO aud f>fiyax/10 to cach p-isotope abundance in Figure 6 (the standard model with 3.5 ranging between 0.5 aud 0.7).," We thus apply the condition $f > 10^5$ and $f >
f_\mathrm{max}/10$ to each p-isotope abundance in Figure 6 (the standard model with $Y_\mathrm{e, 3}$ ranging between 0.5 and 0.7)."
 The p-isotopes that satisfy the above condition are isted in the last colui of Table Ll., The p-isotopes that satisfy the above condition are listed in the last column of Table 4.
 According to recent wdrodvuamiuc studies (Fischeretal.2010:ITüdepohetal.2010).. the maxinuun Y. in the neutrino-driven outflows is ~0.6.," According to recent hydrodynamic studies \citep{Fisc2010, Hued2010}, the maximum $Y_\mathrm{e}$ in the neutrino-driven outflows is $\sim 0.6$."
 Therefore. the p-isotopes that satisfv 1ο above condition only with 355>0.6 are indicate wo c[ves.," Therefore, the p-isotopes that satisfy the above condition only with $Y_\mathrm{e, 3} > 0.6$ are indicated by “[yes]”."
 This implies that the mp-process in core-collapse superuovae is the possible astrophysical origin of 1ο light p-uuclei up to A=108., This implies that the $\nu$ p-process in core-collapse supernovae is the possible astrophysical origin of the light p-nuclei up to $A = 108$.
 In principle. however. je Ép-process can account for the origin of the heavy »iotopes up to A=152 as well. if Y55z0.65 Figure 6) is achieved in the ueutime-driven outflows.," In principle, however, the $\nu
$ p-process can account for the origin of the heavy p-isotopes up to $A
= 152$ as well, if $Y_\mathrm{e, 3} \approx 0.65$ (Figure 6) is achieved in the neutrino-driven outflows."
 If js ds tue. a reasonable combination of the astroplivsica sources considered here can explain all the origins of 1ο solar p-isotopes.," If this is true, a reasonable combination of the astrophysical sources considered here can explain all the origins of the solar p-isotopes."
 It should be noted that most of 1e ΙΑΝ production factors of these heavy. p-uuclei are Zo107., It should be noted that most of the maximum production factors of these heavy p-nuclei are $\gtrsim 10^8$.
 This is three orders of magnitude larger zu the above requisite value (f=10?), This is three orders of magnitude larger than the above requisite value $f = 10^5$ ).
 Thus. oulv OL of neutrino-driven ejecta with 35.420.60.—0.65 is enough to account for the origina of these heavy p-imclei;," Thus, only $\sim 0.1\%$ of neutrino-driven ejecta with $Y_\mathrm{e, 3} \approx 0.60-0.65$ is enough to account for the origin of these heavy p-nuclei."
 Future multi«dineusional lvdvodvuamue studies of core-collapse superuovae with full neutrino transport will be of particular importance if such a condition is indeed obtained., Future multi-dimensional hydrodynamic studies of core-collapse supernovae with full neutrino transport will be of particular importance if such a condition is indeed obtained.
 A word of caution for the ποοσπα isotopes ds needed hore., A word of caution for the molybdenum isotopes is needed here.
" The production factors of ""Mo and ?! Mo satisty the above condition only mareiually with 1,4=0.53. 0.51"," The production factors of $^{92}$ Mo and $^{94}$ Mo satisfy the above condition only marginally with $Y_\mathrm{e, 3} = 0.53-0.54$ ."
 The future measurements of the nuclear masses of “Zr and “Nb wüight in part eme this problem as discussed in 5.3., The future measurements of the nuclear masses of $^{82}$ Zr and $^{83}$ Nb might in part cure this problem as discussed in 5.3.
" This is rather serious for the origin of ?Mo that can be produced only by the mp- while ""Mo can be explained by the p-QSE."," This is rather serious for the origin of $^{94}$ Mo that can be produced only by the $\nu$ p-process, while $^{92}$ Mo can be explained by the p-QSE."
 iskeretal.(2009). couclided that the ratio ??Mo/?! Mo is about 5 times simaller than the solar value. when applving the proton separation enerev of Rl in Weberetal. (2008)...," \citet{Fisk2009} concluded that the ratio $^{92}$ $^{94}$ Mo is about 5 times smaller than the solar value, when applying the proton separation energy of $^{93}$ Rh in \citet{Webe2008}. ."
" This might implies that ""Mo las another origin. presumably the p-QSE."," This might implies that $^{92}$ Mo has another origin, presumably the p-QSE."
 We however obtain a reasonable ratio with our standard model (sec. e.g.. thebottom panel of Figure 22) and πας other cases (see the 35.54<0.55 rauge in Figure 6).," We however obtain a reasonable ratio with our standard model (see, e.g., thebottom panel of Figure 22) and many other cases (see the $Y_\mathrm{e, 3} \le 0.55$ range in Figure 6)."
" This is due to the significant vole of “Rute.p) Te that competes with ""2 fup.5)? Rb in our cases."," This is due to the significant role of$^{92}$ $(n, p)^{92}$ Tc that competes with $^{92}$ $(p, \gamma)^{93}$ Rh in our cases."
" This is a consequence of the values of A, in the present cases being about a factor of three ligher than those iu Pructetal. (2006)...", This is a consequence of the values of $\Delta_\mathrm{n}$ in the present cases being about a factor of three higher than those in \citet{Prue2006}. .
 This suggests that ??Mo/2? ο is highly sensitive to the details of supernova dynamics., This suggests that $^{92}$ $^{94}$ Mo is highly sensitive to the details of supernova dynamics.
 We oeinvestigated the effects of uncertainties in supernova dynamics as well as in nuclear data iuputs ou the rp-process iu the neutrine-driven outflows of core-collapse supernovae., We investigated the effects of uncertainties in supernova dynamics as well as in nuclear data inputs on the $\nu$ p-process in the neutrino-driven outflows of core-collapse supernovae.
" The former includes the winel-termination radius r4 (or temperature Ty). neutrino huuinositv £,. neutron-star mass AL... and clectron fraction νο (or Vig. at To9 and 3. respectively)."," The former includes the wind-termination radius $r_\mathrm{wt}$ (or temperature$T_\mathrm{wt}$ ), neutrino luminosity $L_\nu$ , neutron-star mass $M_\mathrm{ns}$ , and electron fraction $Y_\mathrm{e, 9}$ (or $Y_\mathrm{e, 3}$ , at $T_9 = 9$ and 3, respectively)."
 The latter includes the reactions relevant to the breakout from the pp-chain region (A« 12). the (i.p) reactions on heavy nucleà (Z2 56). aud the unclear masses (40xZ< 50) on the rp-process pathway.," The latter includes the reactions relevant to the breakout from the pp-chain region $A < 12$ ), the $(n, p)$ reactions on heavy nuclei $Z \ge
56$ ), and the nuclear masses $40 \le Z \le 50$ ) on the $\nu$ p-process pathway."
 Our result, Our result
saturation of the dynamo itself (0.8.?).. a saturation of ιο Gilling factor of active regions ou the stellar surface 7).. or a centrifueal stripping of the corona. caused by io high rotation rates (δει,"saturation of the dynamo itself \citep[e.g.][]{vilh84}, a saturation of the filling factor of active regions on the stellar surface \citep{vilh84}, or a centrifugal stripping of the corona caused by the high rotation rates \citep{jard99}."
 However. ouce saturation ceurs the N-rav enmudssion becomes a function of ouly re bolometric buuinositv (7).. or effectively the mass. ‘olor or radius of the mai-sequence star.," However, once saturation occurs the X-ray emission becomes a function of only the bolometric luminosity \citep{pizz03}, or effectively the mass, color or radius of the main-sequence star."
 Iu the non-saturated reeiuec. the two infiuences ou ιο efficiency of the maguetie dviauuo were combined * 7) into a single parameter. the Rossby nunber. Ro=PosfrT. the ratio of the stellar rotation period. D. AUC the mass-dependent convective turnover tine. T.," In the non-saturated regime, the two influences on the efficiency of the magnetic dynamo were combined by \citet{noye84} into a single parameter, the Rossby number, $R_0 = P_{rot} / \tau$, the ratio of the stellar rotation period, $P_{rot}$ , and the mass-dependent convective turnover time, $\tau$."
" This quantity has proven to be au effective parameter of the stellar magnetic dynamo. increasius toward lower masses with the efficiency of the dynamo (eg.22??).,"," This quantity has proven to be an effective parameter of the stellar magnetic dynamo, increasing toward lower masses with the efficiency of the dynamo \citep[e.g.][]{mice84,magg87,step94,rand00}."
 Despite this work there is vet to be a satisfactory dynamo theory that can explain both the solar dvnamo aud that of rapidly rotating stars (c.e.77). and the continued. lack of a sufficicutly large aud unbiased sample has no doubt contributed to this.," Despite this work there is yet to be a satisfactory dynamo theory that can explain both the solar dynamo and that of rapidly rotating stars \citep[e.g.][]{weis05,bran11} and the continued lack of a sufficiently large and unbiased sample has no doubt contributed to this."
 The paucity of stellar samples with which to study the rotationactivity relationship has mainly been due to the difficulty of incasuring accurate stellar rotation periods. which require iultiple deep observations over long baselines.," The paucity of stellar samples with which to study the rotation–activity relationship has mainly been due to the difficulty of measuring accurate stellar rotation periods, which require multiple deep observations over long baselines."
 This has led to the use of projected rotational velocities as a substitute. which are influenced by the nucertaitics of estimated stellar radi and uukuowu inclination angles.," This has led to the use of projected rotational velocities as a substitute, which are influenced by the uncertainties of estimated stellar radii and unknown inclination angles."
 The recent increase i measured rotation periods (c.g.7.increasedthenumberofPleiadesstarswithmeasuredperiodsbyafactoroffive} for uauv thousands of stars in open clusters of known age is overcoming this problem and it is likely that we will iencetforth be limited by the availability of deep X-ray observations for sucht stars., The recent increase in measured rotation periods \citep[e.g.][increased the number of Pleiades stars with measured periods by a factor of five]{hart10} for many thousands of stars in open clusters of known age is overcoming this problem and it is likely that we will henceforth be limited by the availability of deep X-ray observations for such stars.
 Iu this work. we combine new ileasurements of photometric rotation periods for a large number of field aud cluster stars with archival X-ray observatious o produce the largest existing sample of stars with photometric rotation periods and X-ray hunuinosities (Section 2).," In this work, we combine new measurements of photometric rotation periods for a large number of field and cluster stars with archival X-ray observations to produce the largest existing sample of stars with photometric rotation periods and X-ray luminosities (Section 2)."
 This sample is then used in Section 3 to study aud characterize the rotation - activity relationship in detail and to probe the stellar magnetic dynamo responsible for it., This sample is then used in Section 3 to study and characterize the rotation - activity relationship in detail and to probe the stellar magnetic dynamo responsible for it.
 This allows us iu Section | to trace out the N-rav evolution of low-mass stars as a fiction of rotation period. which is a good proxy for age.," This allows us in Section 4 to trace out the X-ray evolution of low-mass stars as a function of rotation period, which is a good proxy for age."
 Finally. in Section 5 this sample is used to derive a new eupirical mcasure of the mass-dependent couvective turnover tine.," Finally, in Section 5 this sample is used to derive a new empirical measure of the mass-dependent convective turnover time."
 To study the relationship between rotation aud activity a sample was compiled from the literature by searching for stars with measurements of both rotation periods and N-ray luminosities., To study the relationship between rotation and activity a sample was compiled from the literature by searching for stars with measurements of both rotation periods and X-ray luminosities.
 Only photometrically-detcruined rotation periods were inchided. discarding all rotation velocity mnieasuremeuts and upper lnits. aud only stars with significant ταν detectious were used. discarding all sources with ouly wpper lanits.," Only photometrically-determined rotation periods were included, discarding all rotation velocity measurements and upper limits, and only stars with significant X-ray detections were used, discarding all sources with only upper limits."
 This choice reduces the sample size available and also has the potential to iuftroduce an X-aav Dmuuinositv bias in our results., This choice reduces the sample size available and also has the potential to introduce an X-ray luminosity bias in our results.
 However this greatlv siupli&es the following analvsis. particularly in the ποτ of the large variety of sources used to compile this sample. the different techniques used to calculate upper hits bv different authors. aud the poteutial inconipleteuesses in upper Its present iu cach sample.," However this greatly simplifies the following analysis, particularly in the light of the large variety of sources used to compile this sample, the different techniques used to calculate upper limits by different authors, and the potential incompletenesses in upper limits present in each sample."
 The inherent biases that will exist in this sample will be discussed aud addressed later., The inherent biases that will exist in this sample will be discussed and addressed later.
 The recent study of the activityrotation relation by 7) provided the starting point for the catalog., The recent study of the activity–rotation relation by \citet{pizz03} provided the starting point for the catalog.
 From their work 102 cluster stars (excluding Pleiades 1ieuibers. which were compiled separately) and 17 field stars were used. excluding all sources with upper limits. as well as a uuniber of stars whose rotation periods were inferred indirectly from chromospheric activity levels (with the exception of a Centauri D. for which we use the rotation period aud mean X-ray Iunuinositv presented by ?))).," From their work 102 cluster stars (excluding Pleiades members, which were compiled separately) and 47 field stars were used, excluding all sources with upper limits, as well as a number of stars whose rotation periods were inferred indirectly from chromospheric activity levels (with the exception of $\alpha$ Centauri B, for which we use the rotation period and mean X-ray luminosity presented by \citet{dewa10}) )."
 The majority of these sources are C and Is stars with the AL stars confined to the field star παπηρ]ο because of the lack of rotational periods available for low mass stars in clusters at the time., The majority of these sources are G and K stars with the M-type stars confined to the field star sample because of the lack of rotational periods available for low mass stars in clusters at the time.
 A further 28 IEvades micuibers were introduced by cross-natching recent rotation periods from ?) with N-vav luminosities from?) aud ?).., A further 28 Hyades members were introduced by cross-matching recent rotation periods from \citet{delo11} with X-ray luminosities from \citet{ster94} and \citet{ster95}.
 The first expansion of the sample was based ou the recent nieasureineut of photometric rotation periods bv 7) for Pleiades members based on the menboership list of ?).., The first expansion of the sample was based on the recent measurement of photometric rotation periods by \citet{hart10} for Pleiades members based on the membership list of \citet{stau07}. .
" These were crossanatehed with N-rayv flux mcasurements of Pleiades stars from the audROSAT observations from ?).. 2).. 2). and. ?).. aud observations frou ο). again discarding all upper Πατ»,"," These were cross-matched with X-ray flux measurements of Pleiades stars from the and observations from \citet{mice90}, \citet{stau94}, \citet{mice96}, and \citet{mice99}, and observations from \citet{brig03}, again discarding all upper limits."
 Where multiple measirements of the X-rav flux exist for a single source. that with the lowest ractional uncertaüntv was used.," Where multiple measurements of the X-ray flux exist for a single source, that with the lowest fractional uncertainty was used."
 Crossanatehiug 391 A-rav sources with the 383 sources with photometric verqods lead to a sample of 116 Pleiades members with voth N-ray luminosities and rotation periods. the largest such sample for a single cluster.," Cross-matching 391 X-ray sources with the 383 sources with photometric periods lead to a sample of 146 Pleiades members with both X-ray luminosities and rotation periods, the largest such sample for a single cluster."
 This sample was then complemented with 83 stars ποια the open clusters NGC 2516 and NCC 2517using. respectively. rotation periods from 7). aud ?).. aud N-rav fluxes from?) and ?)..," This sample was then complemented with 83 stars from the open clusters NGC 2516 and NGC 2547using, respectively, rotation periods from \citet{irwi07} and \citet{irwi08}, and X-ray fluxes from \citet{pill06} and \citet{jeff11}."
 An additional 20 stars were added roni the open cluster Pracsepe using rotation periods tour 2) and ?) and X-ray luminosities from 7) and 7)., An additional 20 stars were added from the open cluster Praesepe using rotation periods from \citet{delo11} and \citet{scho11} and X-ray luminosities from \citet{rand95} and \citet{fran03}.
 The sample was then further extended using data from rotation period surveys of field stars combined with X-rav fluxes from theROSAT Al-Sky Survey., The sample was then further extended using data from rotation period surveys of field stars combined with X-ray fluxes from the All-Sky Survey.
 This included. 218 stars from ?).. 23 stars from ?).. ο stars Toni 7).. 8 stars from 2). aud 79 stars were added from he compilation of ?)..," This included 218 stars from \citet{hart11}, 23 stars from \citet{kira07}, 6 stars from \citet{xing07}, 8 stars from \citet{bouv97} and 79 stars were added from the compilation of \citet{mama08}."
 This sample also inclides the Sin using the values of log Lyρω=6.21 (2) and P426.09 davs (?).., This sample also includes the Sun using the values of log $L_X / L_{bol} = -6.24$ \citep{judg03} and $P_{rot} = 26.09$ days \citep{dona96}.
 Finally. rotation periods cor 65 stars were taken from observations as part of the FEPS (Formation aud Evolution of Planetary Systems. sce Appendix A} program.," Finally, rotation periods for 65 stars were taken from observations as part of the FEPS (Formation and Evolution of Planetary Systems, see Appendix A) program."
 πο classical T-Tinni stars with Πα endüssion. signaling the presence of accretion and therefore a cmeunistella disk. were excluded because of the coniplicatious induced by X-ray. cussion frou accretion and disk-locking ou the rotation period.," Known classical T-Tauri stars with $\alpha$ emission, signaling the presence of accretion and therefore a circumstellar disk, were excluded because of the complications induced by X-ray emission from accretion and disk-locking on the rotation period."
 Pre-MS stars (specifically those with ages ϱ1 Myys) werealso excluded because of potential differences im their internal structure as a function of either mass or cffective temperature., Pre-MS stars (specifically those with ages $<$ 10 Myrs) werealso excluded because of potential differences in their internal structure as a function of either mass or effective temperature.
 Using the catalog ofX-ray variable sources presented by ?).. we removed all ROSAT sources that liad been observed to flare during the observation to lesseu," Using the catalog ofX-ray variable sources presented by \citet{fuhr03}, , we removed all ROSAT sources that had been observed to flare during the observation to lessen"
"For linear perturbations, on large scales, Eq.(1) leads to — where the ó,,,,,0y and 6, are the relative overdensities of PBHs, Poisson fluctuations and radiation, respectively.","For linear perturbations, on large scales, Eq.(1) leads to = where the $\delta_{\pbh}$ $\delta_p$ and $\delta_r$ are the relative overdensities of PBHs, Poisson fluctuations and radiation, respectively."
" Since 6, in Eq.(1)is observableand constant, one would conclude that the quantity — =0 is gauge-invariant (4)and conserved."," Since $\delta_p$ in Eq.(1)is observable constant, one would conclude that the quantity - = _p is gauge-invariant and conserved."
" Indeed this is the entropy per PBH, which should remain constant as long as the universe expands adiabatically (e.g. see Mukhanov 1992)."," Indeed this is the entropy per PBH, which should remain constant as long as the universe expands adiabatically (e.g. see Mukhanov 1992)."
" The associated perturbations, generated in this way are isocurvature(or entropy) perturbations, as the curvature at large scales is not (immediately) affected by the formation of compact objects at small scale."," The associated perturbations, generated in this way are isocurvature(or entropy) perturbations, as the curvature at large scales is not (immediately) affected by the formation of compact objects at small scale."
" As we are assuming that PBHs are the present day Cold Dark Matter (CDM), the overdensity of CDM is given by Tiso(K)S(k), where T44(k) and T;,o(k)οι(8) are the transfer functions for adiabatic and isocurvature perturbations respectively."," As we are assuming that PBHs are the present day Cold Dark Matter (CDM), the overdensity of CDM is given by (k) S(k), where $T_{ad}(k)$ and $T_{iso}(k)$ are the transfer functions for adiabatic and isocurvature perturbations respectively."
 For the following analysis we will use the analytical fits quoted in Bardeen 1986 to the transfer functions., For the following analysis we will use the analytical fits quoted in Bardeen 1986 to the transfer functions.
 Eq. (, Eq. (
"5) leads to the following power spectrum In this expression,P;aaq(k)=Ak” with nm~ is the adiabatic power spectrum which is produced through1 inflation (or an alternative method of generating scale-invariant adiabatic perturbations), while P, is given in Eq.(2).","5) leads to the following power spectrum In this $P_{i,ad}(k) = A\,
k^n$ with $n\simeq 1$ is the adiabatic power spectrum which is produced through inflation (or an alternative method of generating scale-invariant adiabatic perturbations), while $P_p$ is given in Eq.(2)."
" One can easily see that the isocurvature term on the RHS of Eq.(2) contributes a constant to the power spectrum as both P, and Tiso(k) = Zeq)Heq(T) are independent of k (e.g. Peacock 1998).", One can easily see that the isocurvature term on the RHS of Eq.(2) contributes a constant to the power spectrum as both $P_p$ and (k) = ) are independent of $k$ (e.g. Peacock 1998).
 Note that this is the simple linear growth due to gravitational clustering which is the same for adiabatic fluctuation., Note that this is the simple linear growth due to gravitational clustering which is the same for adiabatic fluctuation.
" Since the power spectrum of adiabatic fluctuations decays as k? at small scales, one expects to see the signature of this Poisson noise at large k’s. Combining Eqs. ("," Since the power spectrum of adiabatic fluctuations decays as $k^{-3}$ at small scales, one expects to see the signature of this Poisson noise at large $k$ 's. Combining Eqs. ("
"2),(6) and (7) gives the power offset = 4,63 (Cei) (0... 9) Mpc)*(8) which is also a lower bound on the matter linear power spectrum.","2),(6) and (7) gives the power offset = 4.63 ( ) h^5) )^3 which is also a lower bound on the matter linear power spectrum."
 Fig.(1) shows the linear power spectrum for different masses of the PBHs., Fig.(1) shows the linear power spectrum for different masses of the PBHs.
 We see the Poisson plateau (Eq., We see the Poisson plateau (Eq.
 8) at large k's which drops with decreasing mass., 8) at large k's which drops with decreasing mass.
 The impact of this plateau on the forest power spectrum is discussed in the next section., The impact of this plateau on the forest power spectrum is discussed in the next section.
the Milky Was.,the Milky Way.
 Scanapieco. Ferrara Broadhurst (2000) confirm the above through detailed clynamical modeling of the photo evaporation and ram striping of gas from dSph's due to galactic winds and fountains.," Scanapieco, Ferrara Broadhurst (2000) confirm the above through detailed dynamical modeling of the photo evaporation and ram striping of gas from dSph's due to galactic winds and fountains."
 All this only stresses the fact that these cdeceivingly simple galactic svstems are subject to complex processes which make it cillieult to construct physical models to ceseribe their evolution., All this only stresses the fact that these deceivingly simple galactic systems are subject to complex processes which make it difficult to construct physical models to describe their evolution.
 In our present study we take the SEItUs derived. by LGC (henceeforth. SEIgea) as external constraints on our chemical evolution models. and. hence. obtain interesting restrictions on the time structure anc magnitude of the eas accretion history of Carina. Ursa. Minor. Leo |. and Leo Le," In our present study we take the SFR's derived by HGC (henceforth $SFR_{HGV}$ ) as external constraints on our chemical evolution models, and hence obtain interesting restrictions on the time structure and magnitude of the gas accretion history of Carina, Ursa Minor, Leo I and Leo II."
 In a sense. for these galaxies we know part of the answer in advance. and can hence caleulate the energy input produced by the inferred star formation history. and restrict the possible gas accretion and outllow scenarios. thus obtaining valuable information on the nature of the ISM-SER connection.," In a sense, for these galaxies we know part of the answer in advance, and can hence calculate the energy input produced by the inferred star formation history, and restrict the possible gas accretion and outflow scenarios, thus obtaining valuable information on the nature of the ISM-SFR connection."
 The time evolution of the metallicities is hen a prediction of the model. which we can compare with observed. values.," The time evolution of the metallicities is then a prediction of the model, which we can compare with observed values."
 The models we obtain show a variety. of possibilities or the physical evolution of hese svstems. depending on which parameters one varies {κ» ensure gas is retained. until he luminous galaxy is formed.," The models we obtain show a variety of possibilities for the physical evolution of these systems, depending on which parameters one varies to ensure gas is retained until the luminous galaxy is formed."
 Galaxies showing repeated »eriods of star formation. such as Carina and Leo | in our sample. can onlv be expained with the inclusion of a re-accretion of fresh. gas.," Galaxies showing repeated periods of star formation, such as Carina and Leo I in our sample, can only be explained with the inclusion of a re-accretion of fresh gas."
 MὉ obtain predictions on the otal masses. metallicities and. abundance ratios of the ejected material. as a result. of wing carefully traced. the ohvsies of the gas content. the clilferent SN vields. ancl the inal results of the outflows.," We obtain predictions on the total masses, metallicities and abundance ratios of the ejected material, as a result of having carefully traced the physics of the gas content, the different SN yields, and the final results of the outflows."
 A simple physical criterion is also proposed. as relevant to discriminating dSph galaxies subject to extended and repeated star formation. from those susceptible only to a single burst of activity.," A simple physical criterion is also proposed as relevant to discriminating dSph galaxies subject to extended and repeated star formation, from those susceptible only to a single burst of activity."
 The plan of our paper is as follows: Section 2 presents the details of the enrichment and eas dynamics mocel. with the results once the inferred SER()'s have been introduced as Constraints. presented in Section 3.," The plan of our paper is as follows: Section 2 presents the details of the enrichment and gas dynamics model, with the results once the inferred SFR(t)'s have been introduced as constraints, presented in Section 3."
 Finally. a discussion of our results is given in Section d and Section 5r states our conclusions.," Finally, a discussion of our results is given in Section 4 and Section 5 states our conclusions."
 As mentioned in the introduction. with the possible exception of Sculptor. all attempts at detecting the presence of gas in dSph's have vielclecl only null results.," As mentioned in the introduction, with the possible exception of Sculptor, all attempts at detecting the presence of gas in dSph's have yielded only null results."
 Et is therefore reasonable to assume that the heating and dynamical cllects of star formation have powered winds which resulted in the loss of gas in these systems., It is therefore reasonable to assume that the heating and dynamical effects of star formation have powered winds which resulted in the loss of gas in these systems.
 We shall assume that only SNac type Land LE are responsible for these heating and dynamical processes. and calculate the appearance of galactic wines in dSph's accordingly.," We shall assume that only SNae type I and II are responsible for these heating and dynamical processes, and calculate the appearance of galactic winds in dSph's accordingly."
 “Phe criterion for the establishment of a wind in essence derives from a comparison of the thermal energy of the eas and its gravitational binding energy., The criterion for the establishment of a wind in essence derives from a comparison of the thermal energy of the gas and its gravitational binding energy.
 Lt is hence the structure of the dark matter halo which fixes the boundary. conditions on the problem., It is hence the structure of the dark matter halo which fixes the boundary conditions on the problem.
 In the following sub-section we describe the details of the dark matter haloes used. and the criterion used to identify the formation of a wind.," In the following sub-section we describe the details of the dark matter haloes used, and the criterion used to identify the formation of a wind."
 We have assumed that a ονο spheroical galaxy is a system mace initially of a non-barvonic cark matter halo ancl a xwvonic gas spheroid., We have assumed that a dwarf spheroidal galaxy is a system made initially of a non-baryonic dark matter halo and a baryonic gas spheroid.
 Direct stuclies of rotation curves in dwarf galaxies have shown the density. profiles οἱ these systems to be well described by a constant density core. ollowed by an isothermal region out to the limit of the observations (e.g. Burkert 1995).," Direct studies of rotation curves in dwarf galaxies have shown the density profiles of these systems to be well described by a constant density core, followed by an isothermal region out to the limit of the observations (e.g. Burkert 1995)."
 Observations in low surface xightness galaxies have shown the same results (e.g. de Blok MeCGaugh 1997). and indeed the pattern appears to extend o the dark components of clusters of galaxies imaged in X ravs (e.g. Firmani ct al.," Observations in low surface brightness galaxies have shown the same results (e.g. de Blok McGaugh 1997), and indeed the pattern appears to extend to the dark components of clusters of galaxies imaged in X rays (e.g. Firmani et al."
 2000)., 2000).
 It has also been shown that jügh surface brightness galaxies are also consistent with this yalo structure (Hernandez Gilmore 19982). which it hence seers reasonable to assume as universal.," It has also been shown that high surface brightness galaxies are also consistent with this halo structure (Hernandez Gilmore 1998a), which it hence seems reasonable to assume as universal."
 For the svstems we are treating here. the details of the dark halo bevond the core radius are largely unimportant. as their presence in the halo of the much larger Milkv. Was implies the existence of a tida racius for these galaxies. bevond which the tidal field of our Galaxy tears olf material.," For the systems we are treating here, the details of the dark halo beyond the core radius are largely unimportant, as their presence in the halo of the much larger Milky Way implies the existence of a tidal radius for these galaxies, beyond which the tidal field of our Galaxy tears off material."
 It can be shown that these tida radii are in fact very similar to the core radii of the visible dSph's., It can be shown that these tidal radii are in fact very similar to the core radii of the visible dSph's.
 We therefore take a dark matter (DAL) distribution represented by a constant density out to the tidal radii of each dSph. followed. by an exponential cut-oll starting a core radius (Llernancez Cilmore 1998b).," We therefore take a dark matter (DM) distribution represented by a constant density out to the tidal radii of each dSph, followed by an exponential cut-off starting at core radius (Hernandez Gilmore 1998b)."
" where Roo=fox. A, is the tical radius. and fou is an input parameter in cach model. varving which with respect to unity we can asses the dependence of our results on the details of the dark halo within which the ealaxies are embedded."," = where $R_{core} = f_{DM} R_{t}$, $R_t$ is the tidal radius and $f_{DM}$ is an input parameter in each model, varying which with respect to unity we can asses the dependence of our results on the details of the dark halo within which the galaxies are embedded."
 Indeed. Odenkirchen ct al. (," Indeed, Odenkirchen et al. ("
2001) perform a new survey of Draco using the Sloan Digital Sky Survey. and conclude that the true core radius of the light distribution is larger than previous estimates showed. with no evidence of any tidal features.,"2001) perform a new survey of Draco using the Sloan Digital Sky Survey, and conclude that the true core radius of the light distribution is larger than previous estimates showed, with no evidence of any tidal features."
 This shows that values of tidal radii found presently in the literature could be lower limits in other cases as well., This shows that values of tidal radii found presently in the literature could be lower limits in other cases as well.
 Also. present day. values for this parameter should. be considered: lower limits. for the corresponding time averaged quantities. since the tidal field of the MW could. well have reduced the dark haloes of dSph over time.," Also, present day values for this parameter should be considered lower limits for the corresponding time averaged quantities, since the tidal field of the MW could well have reduced the dark haloes of dSph over time."
 Ehe details of the density cut bexond Fs. indeed. the presence of any dark matter bevond {ων only mareinally alfect our results.," The details of the density cut beyond $R_{core}$, indeed, the presence of any dark matter beyond $R_{core}$ only marginally affect our results."
the LIL regions in M. 101 (Ixennicutt. Carnett 1996).,the HII regions in M 101 (Kennicutt Garnett 1996).
 A single line logarithmic gradient is found., A single line logarithmic gradient is found.
 This might actually constitute a purely observational way to quantify ealactic disc abundance gradients without the need to rely on theoretical photoionization mocels., This might actually constitute a purely observational way to quantify galactic disc abundance gradients without the need to rely on theoretical photoionization models.
 We have performed a new empirical calibration of nebular abundances using the sulphur abundance. parameter σου., We have performed a new empirical calibration of nebular abundances using the sulphur abundance parameter $S_{23}$.
 ‘This calibration is an alternative to the commonly used one based on the strong optical oxvgen lines and presents several advantages., This calibration is an alternative to the commonly used one based on the strong optical oxygen lines and presents several advantages.
 From the observational point of view. the lines are easily observable. both in. low and high. metallicity. regions. and less alfected by reddening.," From the observational point of view, the lines are easily observable, both in low and high metallicity regions, and less affected by reddening."
 Furthermore. their intensities can be measured relative to nearby. hydrogen recombination lines thus minimizing any cllects due to uncertainties in Εαν calibration.," Furthermore, their intensities can be measured relative to nearby hydrogen recombination lines thus minimizing any effects due to uncertainties in flux calibration."
 On the theoretical side. their contribution to the cooling of the nebula becomes important at electron temperatures lower than in the case of the traditional Os; (previously called. 1253) and therefore its relation with oxvgen abundance remains single-valued up to metallicities close to solar.," On the theoretical side, their contribution to the cooling of the nebula becomes important at electron temperatures lower than in the case of the traditional $O_{23}$ (previously called $R_{23}$ ) and therefore its relation with oxygen abundance remains single-valued up to metallicities close to solar."
 Also. the fact that $2 ds less dependent than Qo; on ionization parameter reduces the seatter in the relation.," Also, the fact that $S_{23}$ is less dependent than $O_{23}$ on ionization parameter reduces the scatter in the relation."
 The application of this new metallicity calibration can provide more accurate abundance determinations for objects with logO»; between 0.5 and. 1.2. oxygen abundances between 12]|log(O/Il)2 7.20 (c 0.02 times solar) and 121log(O/LI) = S.80 (2 0.75 times solar).," The application of this new metallicity calibration can provide more accurate abundance determinations for objects with $O_{23}$ between 0.5 and 1.2, oxygen abundances between 12+log(O/H)= 7.20 $\simeq$ 0.02 times solar) and 12+log(O/H) = 8.80 $\simeq$ 0.75 times solar)."
 This is the range of metallicities found in LIL galaxies and HEIL regions in irregular galaxies and outer galactic clises., This is the range of metallicities found in HII galaxies and HII regions in irregular galaxies and outer galactic discs.
 Regarding LIL regions of higher metallicity. the composedparameter SoOo; ," Regarding HII regions of higher metallicity, the composedparameter $S_{23}/O_{23}$ "
"Zp3. 3.:3«10? - LasLo? 7. the averaged value is 1.1«1019 3,","ZP3, $3.3\times10^{9}$ - $4.0\times10^{9}$ $^{-3}$, the averaged value is $1.4\times10^{10}$ $^{-3}$."
 CGeuerallv. the source regiou with plasma density as up to 5.5«Lott P is always located very. close to the base of solar corona where the height frou the solar photosphere is oulv several thousauds kilometers. while the source region with plasma deusity of about «1019} 3 is located near the bottom of the solar corona.," Generally, the source region with plasma density as up to $5.5\times10^{11}$ $^{-3}$ is always located very close to the base of solar corona where the height from the solar photosphere is only several thousands kilometers, while the source region with plasma density of about $\times10^{10}$ $^{-3}$ is located near the bottom of the solar corona."
 However. it should be differeut around the active regions.o especially arouud the flaring regions.," However, it should be different around the active regions, especially around the flaring regions."
 The X-ray observatious indicate that the plagma densities around the flaviug core region are in the range of 1091012 «mn7 and their heights can be iu several decades of thousauds kin above the solar photosphere (Olivamia Shibata. 1998).," The X-ray observations indicate that the plasma densities around the flaring core region are in the range of $10^{9} - 10^{11}$ $^{-3}$, and their heights can be in several decades of thousands km above the solar photosphere (Ohyama Shibata, 1998)."
 Oue of the crucial and most difficult problem in solar plysics is to determine the coronal magnetic field confidently., One of the crucial and most difficult problem in solar physics is to determine the coronal magnetic field confidently.
 There are many publications which preseut tle estimations of the coronal magnetic field by using solar radio observations (Maun. Iaulicky. Motschinann. 1987: (οσοκα. 1998: Tang Nakajima. 2002: Thang 2008. ote).," There are many publications which present the estimations of the coronal magnetic field by using solar radio observations (Mann, Karlicky, Motschmann, 1987; Gelfreikh, 1998; Huang Nakajima, 2002; Huang 2008, etc.)."
 Recent observations of microwave bursts with fine structures open up a new possibilities for determining the coronal magnetic Seld (Ikarliekwv Jiricka. 1995: Lenedey et al. 2001. ete).," Recent observations of microwave bursts with fine structures open up a new possibilities for determining the coronal magnetic field (Karlicky Jiricka, 1995; Lenedev et al, 2001, etc)."
 The ZP structure is oue of most important microwave fine structures which can be used to diaguose magnetic field streneth in the coronal source regions. although the results depeud on the theoretical models.," The ZP structure is one of most important microwave fine structures which can be used to diagnose magnetic field strength in the coronal source regions, although the results depend on the theoretical models."
 Different ZP model will deduce differeut values of magnetic field iu the ZP source region., Different ZP model will deduce different values of magnetic field in the ZP source region.
 Practically. it is always difficult to verdict which model is the best one fitted to observations.," Practically, it is always difficult to verdict which model is the best one fitted to observations."
 Possibly. from the estimations of the magnetic field streneths from the ZP structures. we could eet a considerable restriction for the theoretical models. (," Possibly, from the estimations of the magnetic field strengths from the ZP structures, we could get a considerable restriction for the theoretical models. ("
1) DM model indicates that the frequency separation of the adjaceut zebra stripes is just equal to the electron evro-frequency.,1) BM model indicates that the frequency separation of the adjacent zebra stripes is just equal to the electron gyro-frequency.
 From this we may obtain a direct incaswrement of the magnetic field in the coronal ποιαος region: Tere. the unit of B is m Gauss. and f in Wz.," From this we may obtain a direct measurement of the magnetic field in the coronal source region: Here, the unit of $B$ is in Gauss, and $f$ in Hz."
 Substituting the frequency separation of ZPI. ZP2. aud ZP3 iuto the above expression. we nav ect the magnetic field streneth as 2813.06. 21. 25 C aud Ὁ G Ci. respectively.," Substituting the frequency separation of ZP1, ZP2, and ZP3 into the above expression, we may get the magnetic field strength as 28 – 43 G, 21 – 25 G, and 5 – 6 G, respectively."
 Iu this regime. the magnetic field strength is ouly depending ou the frequency separation between the adjaceut zebra stripes. (," In this regime, the magnetic field strength is only depending on the frequency separation between the adjacent zebra stripes. ("
2) From WW ποσο]. we may get the magnetic field streneth in ZP source region: With this relation. the maeuetic field streneth is two times of that estiiated from DM model: 55. 85 Ci. 12. I9. and 10. 11 Ce. corresponding to ZP1. ZP2. aud ZP3. respectively.,"2) From WW model, we may get the magnetic field strength in ZP source region: With this relation, the magnetic field strength is two times of that estimated from BM model: 55 – 85 G, 42 – 49 G, and 10 – 11 G, corresponding to ZP1, ZP2, and ZP3, respectively."
 This regiae is also independent to the inhomogeneous scale height in the source region. (, This regime is also independent to the inhomogeneous scale height in the source region. (
3) From DPR model. we may obtain the measurement of magnetic field streugth in the ZP structure source region.,"3) From DPR model, we may obtain the measurement of magnetic field strength in the ZP structure source region."
" Based oun Equation (3) aud (1). the magnetic field streneth can be derived: Tere. Q is an inhomogeneous factor which is dominated mainly by the scale heights of plasiua density n, and the magnetic field B in the source region."," Based on Equation (3) and (4), the magnetic field strength can be derived: Here, $Q$ is an inhomogeneous factor which is dominated mainly by the scale heights of plasma density $n_{e}$ and the magnetic field $B$ in the source region."
 It can be expressed as:, It can be expressed as:
cohbunon features with the solar corona (e.e..Galeev.Rosner. 1979). and magnetic reconnection iu AGN coronae is a good candidate for the origin of hot electrons(Liu.Mineshige. 2002).,"common features with the solar corona (e.g., 1979), and magnetic reconnection in AGN coronae is a good candidate for the origin of hot electrons, 2002)."
 It is well known that particles are accelerated to nouthermal energies by reconnections in solar flares 1995)., It is well known that particles are accelerated to nonthermal energies by reconnections in solar flares 1995).
" Tere we coustruct a new model of the N/eanunia-axy spectra of Αννα, bv calculating the Couptonizatiou xocess bv hot electrons having both thermal aud rtonthermal couponeuts."," Here we construct a new model of the X/gamma-ray spectra of AGNs, by calculating the Comptonization process by hot electrons having both thermal and nonthermal components."
 We also calculate the CXD spectrum based ou our model with the latest knowledge of the cosmological evolution of the ACN luminosity Muction. and determine the amount aud spectra of ie nonthermal electrons in ACN coronae to explain 1ο MeV background.," We also calculate the CXB spectrum based on our model with the latest knowledge of the cosmological evolution of the AGN luminosity function, and determine the amount and spectrum of the nonthermal electrons in AGN coronae to explain the MeV background."
 We discuss the implied nature of ronthermal electrons in the contest of the reconnection jieatiue scenario of the AGN coronae. comparing our results with those found in the reconnections occurriug in the solar flares and the Earth maguetosphere. (," We discuss the implied nature of nonthermal electrons in the context of the reconnection heating scenario of the AGN coronae, comparing our results with those found in the reconnections occurring in the solar flares and the Earth magnetosphere. ("
1991) and (1993) xeseuted an ACN spectral model that can explain tle MeV backeround spectra by nouthermal relativistic electrons.,1991) and (1993) presented an AGN spectral model that can explain the MeV background spectrum by nonthermal relativistic electrons.
" IHowever. their model ouly cousiclers the ronthermal componcut without a thermal compoucut. and if requires a cut-off of 5,~30 in the ronthermal compoucut. which is difficult to interpret as rev mmcutionced in their paper."," However, their model only considers the nonthermal component without a thermal component, and it requires a cut-off of $\gamma_e \sim 30$ in the nonthermal component, which is difficult to interpret as they mentioned in their paper."
 Our model cousiders both ιο thermal aud nouthermal coronal electrons whose spectra are smoothly connected to cach other. which is a watural extension of the popular ACN spectral models iu re recent Literature.Stecker.Salamon. (," Our model considers both the thermal and nonthermal coronal electrons whose spectra are smoothly connected to each other, which is a natural extension of the popular AGN spectral models in the recent literature., ("
1999) also discussed a possibility that the MeV. backeround is explained by nouthermal tails in ACN spectra. quoting 1ο spectrum of the Calactic stellar-nass black hole candidate Cve X-1.,"1999) also discussed a possibility that the MeV background is explained by nonthermal tails in AGN spectra, quoting the spectrum of the Galactic stellar-mass black hole candidate Cyg X-1."
 However. a plivsical model to explain 1ο nonthermal tail in an ACN spectrum was not xeseuted.," However, a physical model to explain the nonthermal tail in an AGN spectrum was not presented."
 Throughout this paper. we adopt the cosmological λαοτους of (69.ιν04 )2(0.7.0.3.0.7).," Throughout this paper, we adopt the cosmological parameters of $(h_0,\Omega_m,\Omega_\lambda$ )=(0.7,0.3,0.7)."
" The main shape of N-vay ACN spectra is determined bv Conptonization of UV photons ciuitted from optically-thick aceretion disks bv hot olectrous in coronac,", The main shape of X-ray AGN spectra is determined by Comptonization of UV photons emitted from optically-thick accretion disks by hot electrons in coronae.
 As in many previous studies. we consider a simple spherical and wuiform distribution of the coronal electrons.," As in many previous studies, we consider a simple spherical and uniform distribution of the coronal electrons."
 The seed UW photons are injected at the center and then become X-ray photons when they escape the coronal region after Coniptouization., The seed UV photons are injected at the center and then become X-ray photons when they escape the coronal region after Comptonization.
" Iu addition to the hot thermal electrous assumed iu the couventional ANorav spectral models of AGNs. we introduce higher enerev nouthermal electrons in ACN coronae. whose οποιον distribution is a power-law as dN,fdE,xT "," In addition to the hot thermal electrons assumed in the conventional X-ray spectral models of AGNs, we introduce higher energy nonthermal electrons in AGN coronae, whose energy distribution is a power-law as $dN_e/dE_e \propto E_e
^{-\Gamma}$."
"We iutroduce the transition clectron Lorentz factor u. Corresponding to the transition oelectrou cuerey DL,-—ety, Where the electron spectrum NV,ας, has the same value for the thermal aud nonthermal components."," We introduce the transition electron Lorentz factor $\gamma_{\rm{tr}}$, corresponding to the transition electron energy $E_e
= m_e \gamma_{\rm tr}$ where the electron spectrum $dN_e/dE_e$ has the same value for the thermal and nonthermal components."
" This 54 is the lower limit of the Lorentz factor distribution of the uouthermal component. aud hence there are no nonthermal electrons at E,κ σεν."," This $\gamma_{\rm
tr}$ is the lower limit of the Lorentz factor distribution of the nonthermal component, and hence there are no nonthermal electrons at $E_e < m_e \gamma_{\rm tr}$ ."
" We also set an upper bound as 5,=100, although this rardly affects our results if the πακατα photon energy well extends bevoud LO MeV. We set the coronal teiiperature to be ET,=256 τον and assume a blackbody spectrum for UV seed xhotous from a cooler disk with 7;= 10 eV. following he conventional thermal models (c.g. 1991)."," We also set an upper bound as $\gamma_{u} =
10^5$, although this hardly affects our results if the maximum photon energy well extends beyond 10 MeV. We set the coronal temperature to be $kT_e=256$ keV and assume a blackbody spectrum for UV seed photons from a cooler disk with $T_d =
$ 10 eV, following the conventional thermal models (e.g. 1994)."
 The deeree of Comptonization is determined by he optical depth for Thomson scattering. rr.," The degree of Comptonization is determined by the optical depth for Thomson scattering, $\tau_T$ ."
 We fouud hat the spectral photon iudex in the N-rav baud. ay. »ecomies Close to that typically fouud iu observed spectra (ayzm1.9. ee. 1991: 1997: 1998) when we set rr=O21.," We found that the spectral photon index in the X-ray band, $\alpha_X$, becomes close to that typically found in observed spectra $\alpha_X \approx
1.9$, e.g., 1994; 1997; 1998) when we set $\tau_T =
0.24$."
 Therefore we use this value throughout this letter: this value is also simular to those used iu the couventional models., Therefore we use this value throughout this letter; this value is also similar to those used in the conventional models.
 It should be noted that ay is hardly chauged even if we introduce the nouthermal clectrou componoeut with an amount that is necessary to explain the cosmic MeWV backeround., It should be noted that $\alpha_X$ is hardly changed even if we introduce the nonthermal electron component with an amount that is necessary to explain the cosmic MeV background.
 We then trace the Comptonization process using a Monte Carlo method., We then trace the Comptonization process using a Monte Carlo method.
 The calewlation method used here is nmudulv based on that iuPozduiakov.Sobol. (LOTT) aud (1981). but their original forinalizii iu the laboratory frame is not optimized for the ultra-relativistic region.," The calculation method used here is mainly based on that in, (1977) and (1984), but their original formalism in the laboratory frame is not optimized for the ultra-relativistic region."
 To calculate the scattering by high energy nouthermal electrous more efiiciently. we added a new formulation iu the rest frame of relativistic electrons based on (1970).," To calculate the scattering by high energy nonthermal electrons more efficiently, we added a new formulation in the rest frame of relativistic electrons based on (1970)."
 The reflection of X-ray photons by cool optically thick aatter is also an nuportaut feature of Αν X-ray spectra.," The reflection of X-ray photons by cool, optically thick matter is also an important feature of AGN X-ray spectra."
 We calculate this bv using the PEXNRAY model 1995) iu the NSPEC package as done in C03., We calculate this by using the PEXRAV model 1995) in the XSPEC package as done in U03.
 Because of the limitation for the acceptable input spectzuii in PEXRAV. we use a power-law spectrum (ay= 1.9) plus an exponential cutoff at E= 500 keV. which is a good approximation of the Comptonized spectrum only with the thermal clectrous.," Because of the limitation for the acceptable input spectrum in PEXRAV, we use a power-law spectrum $\alpha_X = 1.9$ ) plus an exponential cutoff at $E = $ 500 keV, which is a good approximation of the Comptonized spectrum only with the thermal electrons."
 The newly added nonthermal electrons would change the spectrum significantly only at E2 1 MeV. and hence this troatinent is appropriate for the reflection component which is iuportaut oulv at —1100 keV. We calculate the CNB spectrum by. integrating our AGN spectral model in the redshift and hunuinositv space. using the N-rav ACN luninosity function of U(3.," The newly added nonthermal electrons would change the spectrum significantly only at $E \gtrsim$ 1 MeV, and hence this treatment is appropriate for the reflection component which is important only at $\sim$ 1–100 keV. We calculate the CXB spectrum by integrating our AGN spectral model in the redshift and luminosity space, using the X-ray AGN luminosity function of U03."
" Following the same formulation given in C03. we take iuto account the absorption column density distribution (Ny, function) aud the coutributiou from Comptou-thick AGNs."," Following the same formulation given in U03, we take into account the absorption column density distribution $N_{\rm H}$ function) and the contribution from Compton-thick AGNs."
 We confirm that our main couclusion hardly change if instead we use a dore recent population svuthesis model by (2007). as described below.," We confirm that our main conclusion hardly change if instead we use a more recent population synthesis model by (2007), as described below."
 Figure 1. shows the models of AGN spectra caleulated according to the procedures iu the previous section., Figure \ref{agn} shows the models of AGN spectra calculated according to the procedures in the previous section.
 Tere. we do not take iuto account the reflection component aud 1e absorption effect. to show the pure spectrum of the van'onrptouization.," Here, we do not take into account the reflection component and the absorption effect, to show the pure spectrum of the Comptonization."
 We set P—3.8 and 54=L1 as our standard model (solid line). because we will find that ese values eive the best-fit MeV. backeround spectruuu o the data.," We set $\Gamma=3.8$ and $\gamma_{\rm{tr}}=4.4$ as our standard model (solid line), because we will find that these values give the best-fit MeV background spectrum to the data."
 In this standard model. 3.5% of the total electron energv is carried bythe nouthermal electrons.," In this standard model, $3.5\%$ of the total electron energy is carried bythe nonthermal electrons."
" Toillustrate effects of changing parameters, we also show je spectra with parameters sliehtlv chaneed from those"," Toillustrate effects of changing parameters, we also show the spectra with parameters slightly changed from those"
"magnetic reconnection regions is established iu solar flares (οιοι, Masuda et al.","magnetic reconnection regions is established in solar flares (e.g., Masuda et al."
 1995). reconmection processes lay happen in the chromosphere aud transition region CÀselisvandenu ct al.," 1995), reconnection processes may happen in the chromosphere and transition region (Aschwanden et al."
 2007: Cudiksen and Nordluud 2005a.b). causing subsequent upflows of heated plastua iuto the coronal parts of active region loops.," 2007; Gudiksen and Nordlund 2005a,b), causing subsequent upflows of heated plasma into the coronal parts of active region loops."
 It is too carly to speculate about the details of the generic heating mechamisi of active region loops. before we analyzed comprehcusive iulti5vaveleueth observations of coronal loops such as with ATA and modeled their livdrodyuauic evolution selt-cousisteutlv.," It is too early to speculate about the details of the generic heating mechanism of active region loops, before we analyzed comprehensive multi-wavelength observations of coronal loops such as with AIA and modeled their hydrodynamic evolution self-consistently."
 Our study of the cross-sectional temperature structure of coronal loops using ATA six-filter data leads us to the following conclusions: Future loop studies with ATA are anticipated that determine the thermal loop structure along the loop axis. as well as a function of time. which will provide uuprecedenuted input for hydrodvuamic simulations," Our study of the cross-sectional temperature structure of coronal loops using AIA six-filter data leads us to the following conclusions: Future loop studies with AIA are anticipated that determine the thermal loop structure along the loop axis, as well as a function of time, which will provide unprecedented input for hydrodynamic simulations"
exited modes. and indeed we should not expect a detection of stochastically excited nonradial modes in giants.,"exited modes, and indeed we should not expect a detection of stochastically excited nonradial modes in giants."
 The amplitudes of intrinsically stable stochastically driven racial modes were estimated in the manner of ((1999): where here πω is the noise generation rate injected into à mode through the Ductuating Revnoleds stresses. the expression for which we adopte from Balmforth (1992b) (see also1999).," The amplitudes of intrinsically stable stochastically driven radial modes were estimated in the manner of (1999):, where here $P_Q$ is the noise generation rate injected into a mode through the fluctuating Reynolds stresses, the expression for which we adopted from Balmforth (1992b) (see also."
". The damping rate is yo= =. and 4,=IH2 in our notation."," The damping rate is $\eta=D_{\rm p}/2 I\omega^2=-\gamma$ , and $I_\omega=IR^{-2}$ in our notation."
 Lor radial modes the total energy. dissipation rate Lis η., For radial modes the total energy dissipation rate $D$ is $D_{\rm p}$.
 The linear stability analysis also provides the parameter À. which is the ratio of the relative luminosity to the relative velocity amplitude. computed at the surface (ie. outermost meshpoint) of the star.," The linear stability analysis also provides the parameter $\lambda$, which is the ratio of the relative luminosity to the relative velocity amplitude, computed at the surface (i.e. outermost meshpoint) of the star."
 “Phe bolometric relative luminosity amplitude then becomeseR and from equation (21) we obtain. where 2=Poly and fy13MIP: Ly ds the dimensionless modal inertia. plotted in Fig.," The bolometric relative luminosity amplitude then becomes, and from equation (21) we obtain, where $P=P_QI_{\rm n}$ and $I_{\rm n}=I/3MR^2$; $I_{\rm n}$ is the dimensionless modal inertia plotted in Fig."
 3., 3.
 In the lower panel of Fig., In the lower panel of Fig.
 6. we plot the quantity A?P.," 6, we plot the quantity $\lambda^2P$."
 ALL the quantities plotted in this figure are applicable also to nonracdial modes of low degree., All the quantities plotted in this figure are applicable also to nonradial modes of low degree.
 Llowever. for nonradial modes we have to take into account the damping elfects in the eamode propagation zone.," However, for nonradial modes we have to take into account the damping effects in the g-mode propagation zone."
 With the help of equation (21) and the data eiven in Fig., With the help of equation (21) and the data given in Fig.
" 3 we can evaluate amplitudes for radial modes with I,« 0.", 3 we can evaluate amplitudes for radial modes with $D_{\rm p}<0$ .
 In Fig., In Fig.
" 7. we compare racial-niocle frequencies and amplitudes calculated for AL, with the observational data of a UMa."," 7, we compare radial-mode frequencies and amplitudes calculated for $_\alpha$ with the observational data of $\alpha\,$ UMa."
 Keeping in mind the laree observational errors and the fact tha we have made no effort to adjust. mocel parantCrs ο fit the frequencies. we regard the agreement of frequencies as satisfactory.," Keeping in mind the large observational errors and the fact that we have made no effort to adjust model parameters to fit the frequencies, we regard the agreement of frequencies as satisfactory."
 On the other hand. the disagreement. between the amplitudes is very serious: the observed amplitude at i=1 exceeds the predicted value by three orders of magnitude. ancl the frequency dependence of the amplitudes cliller drasticalls.," On the other hand, the disagreement between the amplitudes is very serious: the observed amplitude at $n=1$ exceeds the predicted value by three orders of magnitude, and the frequency dependence of the amplitudes differ drastically."
 An additional cilliculty is presented by the presence of the two peaks above the acoustic cut-oll frequency., An additional difficulty is presented by the presence of the two peaks above the acoustic cut-off frequency.
 Such high-frequency peaks are observed in the Sun. but with amplitudes much lower than those below the acoustic cut-oll.," Such high-frequency peaks are observed in the Sun, but with amplitudes much lower than those below the acoustic cut-off."
 The two highest-frequceney peaks in à UMa have zumplitudes of about mmmaeg. which are similar to most of the other peaks.," The two highest-frequency peaks in $\alpha\,$ UMa have amplitudes of about mmag, which are similar to most of the other peaks."
 We should stress that the amplitude: estimates in lig., We should stress that the amplitude estimates in Fig.
 7 were obtained using the pulsation modes of a model with an atmosphere based on model € of Vernazza. Averett Loeser (1981).," 7 were obtained using the pulsation modes of a model with an atmosphere based on model C of Vernazza, Avrett Loeser (1981)."
 Phat atmosphere has an acoustic cut-olf frequeney of 32. 4yrllzat the temperature minimum. which is," That atmosphere has an acoustic cut-off frequency of $32.4 \mu$ Hzat the temperature minimum, which is"
magnitude.,magnitude.
 A similar additional scatter is also observed in the rotational WC3N transitions reffie2))., A similar additional scatter is also observed in the rotational $_3$ N transitions \\ref{fig2}) ).
 These cover ouly a sinall amount of excitation. with the highest level. J=7. being just ~12 KI above the eround state.," These cover only a small amount of excitation, with the highest level, $J$ =7, being just $\sim$ K above the ground state."
 Given the velocity scatter im the trausitious of NIT; aud IIC4N. constraining the variation of ji is best done with conservative velocity. aud uncertaüutv estimates which aturally incorporate the observed scatter.," Given the velocity scatter in the transitions of $_3$ and $_3$ N, constraining the variation of $\mu$ is best done with conservative velocity and uncertainty estimates which naturally incorporate the observed scatter."
 The simplest such velocity estimator is obviously the muweighted mean velocity and its standard deviation which. for the NID; ransitious in Table 6. ave &8.9040.2 and for he ΠοΝ rausitions. ave &.58+0.37 +.," The simplest such velocity estimator is obviously the unweighted mean velocity and its standard deviation which, for the $_3$ transitions in Table 6, are $\pm$ $^{-1}$, and for the $_3$ N transitions, are $\pm$ $^{-1}$."
 Of course. he statistical velocity uncertainties quo in Table 6 Or sone transitions are so high that iucluine them in he mean velocity caleulatiou is likely to decrease the reliability of the mean.," Of course, the statistical velocity uncertainties quoted in Table 6 for some transitions are so high that including them in the mean velocity calculation is likely to decrease the reliability of the mean."
" We therefore reject transitions with velocity uncertainties larger than the root-lucall-square (RAIS) velocity variation for cach species. 0.91 and + for NIE, and HIC4N respectively."," We therefore reject transitions with velocity uncertainties larger than the root-mean-square (RMS) velocity variation for each species, 0.91 and $^{-1}$ for $_3$ and $_3$ N respectively."
 With this criterion. the NIT; (10.10) aud Που Ὁς transitions are rejected.," With this criterion, the $_3$ (10,10) and $_3$ N $\leftarrow$ 4 transitions are rejected."
5. Using these final clipped. mean velocities. AV = OSZEO. ss!.," Using these final clipped mean velocities, $\Delta V$ = $\pm$ $^{-1}$."
 Equation 1 then provides our 1-6 constraint on the variation iu p. μμ = O.OSSEO.17) κ *.," Equation \ref{eq:mu} then provides our $\sigma$ constraint on the variation in $\mu$, $\Delta\mu/\mu$ = $\pm$ 0.47) $\times$ $^{-6}$."
 Since the quoted uucertaintv derives cutirely roni the scatter in the individual transition velocities. his should be a reasouablv robust error estimate.," Since the quoted uncertainty derives entirely from the scatter in the individual transition velocities, this should be a reasonably robust error estimate."
 Nevertheless. given that we have used ouly single Caussiau fits to the absorption profiles and that there is some scatter in the oeicdividua transition velocities. we quote our final result as a 23-0 upper lint ou variation in jr at he absorption redshitt of +=0.89. For comparison with laboratory constraints on variatious in yp. aud in the absence of a reliable model or how µ inight be expected to vary with cosmological ine. it Is common. if not well motivated. to assume that any variation is incar in time.," Nevertheless, given that we have used only single Gaussian fits to the absorption profiles and that there is some scatter in the individual transition velocities, we quote our final result as a $\sigma$ upper limit on variation in $\mu$ at the absorption redshift of $z=0.89$, For comparison with laboratory constraints on variations in $\mu$, and in the absence of a reliable model for how $\mu$ might be expected to vary with cosmological time, it is common, if not well motivated, to assume that any variation is linear in time."
 Deuce. our upper luit ou variation in p translates to a 3-0 upper limit on its time variatiou over the past 7.0 Cor.," Hence, our upper limit on variation in $\mu$ translates to a $\sigma$ upper limit on its time variation over the past 7.0 Gyr."
 There is ve another study on the proton-to-clectrou lass ratio 1n he main lens of T1830211., There is yet another study on the proton-to-electron mass ratio in the main lens of 1830–211.
 From a colmparisou of the auumouia inversion lines (Ieukel ct al., From a comparison of the ammonia inversion lines (Henkel et al.
 2008) with its (J.A) = (1.0)< (0.0) rotational trausition. AMeuten et al. (," 2008) with its $J$ $K$ ) = $\leftarrow$ (0,0) rotational transition, Menten et al. ("
2008) fud cousisteucv. at a lo level. of Ap/pe = 1.9<10 ©,"2008) find consistency, at a $\sigma$ level, of $\Delta\mu$ $\mu$ = $\times$ $^{-6}$."
 The streneth of this study is its focus ou lines with different depeudencies on j arius from the same molecular species., The strength of this study is its focus on lines with different dependencies on $\mu$ arising from the same molecular species.
 Ou the other laud. the ratio between the frequencies of the rotational aud inversion lues is —25. thus leading to potentially siguificaut differences in the morphology aud the coveriug factor of the backerouud radio and subnüillimeter coutim£4y.," On the other hand, the ratio between the frequencies of the rotational and inversion lines is $\sim$ 25, thus leading to potentially significant differences in the morphology and the covering factor of the background radio and submillimeter continuum."
. While such systematic differences cannot be quautified ou the basis of a single rotational line. the given uncertaiuty of the resulting Aye /ye value is dominated by the limited signal-to-noise ratio of the rotational line.," While such systematic differences cannot be quantified on the basis of a single rotational line, the given uncertainty of the resulting $\Delta\mu$ $\mu$ value is dominated by the limited signal-to-noise ratio of the rotational line."
" The cousisteney of NIT; aud ουν radial velocities is reniarkable in view of a umber of effects. which ΠΕη, exert a significaut influence on our results."," The consistency of $_3$ and $_3$ N radial velocities is remarkable in view of a number of effects, which might exert a significant influence on our results."
 Ἠανίπο carefully avoided the use of optically thick transitions. these aro (1) time variability of the continuum source. (2) a frequeney depeudent continu morphology. (3) hvperfne structure. (1) chemustry. and (5) iuhomioseucities in temperature aud deusitv inside a cloud of size Z10ppc (Carilli et al.," Having carefully avoided the use of optically thick transitions, these are (1) time variability of the continuum source, (2) a frequency dependent continuum morphology, (3) hyperfine structure, (4) chemistry, and (5) inhomogeneities in temperature and density inside a cloud of size $\ga$ pc (Carilli et al."
 1998). (, 1998). (
1) À time variable contimuni source Way leac to different Lues-of-sielt and thus to different radial velocities at different epochs.,1) A time variable continuum source may lead to different lines-of-sight and thus to different radial velocities at different epochs.
 Not much of his was seen in spite of the existence of monitoring programs (Wikliud Combes 1998: Muller et al., Not much of this was seen in spite of the existence of monitoring programs (Wiklind Combes 1998; Muller et al.
 2006: Muller Cuélliu 2008) or repeated measurements of ammonia lines (Table 6 aud IIenkel et al., 2006; Muller Guéllin 2008) or repeated measurements of ammonia lines (Table 6 and Henkel et al.
 2008). (, 2008). (
2) The NII; and Που transitions used iu this study are mich closer iu frequency than those chosen bx Απρ et al. (,2) The $_3$ and $_3$ N transitions used in this study are much closer in frequency than those chosen by Murphy et al. (
20082) aud Meuteu et al. (,2008a) and Menten et al. (
2008).,2008).
 Also. in 11830211. the source covering factor may be less frequency dependent than in DO218|357. 101.1).," Also, in 1830–211, the source covering factor may be less frequency dependent than in B0218+357 4.1.1)."
 Nevertheless. NIL; aud WC3N frequencies differ by factors of 1.02.5.," Nevertheless, $_3$ and $_3$ N frequencies differ by factors of 1.0–2.5."
 While no direct effect is apparent. it remaius a source of uncertainty which cannot be quantified. (," While no direct effect is apparent, it remains a source of uncertainty which cannot be quantified. ("
3) Iu. particular NIL; velocities would be ereatlv affected by hyperfine (hf) splitting. if strong non-LTE effects would occur.,"3) In particular $_3$ velocities would be greatly affected by hyperfine (hf) splitting, if strong non-LTE effects would occur."
 A main group of if-coniponents ids surrounded by four satellite groups. displaced by about +10 and +20 l," A main group of hf-components is surrounded by four satellite groups, displaced by about $\pm$ 10 and $\pm$ $^{-1}$."
 The üeher the energy above the eround state of an inversion line. the weaker the satellite features relative to the nain eroup (e.g. the satellites accouut for ~2.5% of the total absorption iu the (C7.KN) = (7.7) line).," The higher the energy above the ground state of an inversion line, the weaker the satellite features relative to the main group (e.g., the satellites account for $\sim$ of the total absorption in the $J$ $K$ ) = (7,7) line)."
 The above made comparison of average velocities of the three lowest with the seven higher excited NIL; inversion lines does not show a siguificaut shift in velocity., The above made comparison of average velocities of the three lowest with the seven higher excited $_3$ inversion lines does not show a significant shift in velocity.
 Deviations from LTE intensity ratios of th| various components are onlv expected in the case of a sigmificant population of the nommetastable inversion states (C2 A). which requires extremely veh densities or intense radiation fiekls (Stutziki et al.," Deviations from LTE intensity ratios of the various components are only expected in the case of a significant population of the non-metastable inversion states $J$$>$$K$ ), which requires extremely high densities or intense radiation fields (Stutzki et al."
 19814: Stutzki Winnewisser 1985a.b).," 1984; Stutzki Winnewisser 1985a,b)."
" Non-LTE effects leading to IEC4N velocity shifts as a ""nctiou of rotational quanti umnboer J are also not obvious.", Non-LTE effects leading to $_3$ N velocity shifts as a function of rotational quantum number $J$ are also not obvious.
 In us case. the —3« 2 Line would be the most critical. while he hf-structure of the τς 6 line is far too compact to vield anv significa shifts.," In this case, the $J$ $\leftarrow$ 2 line would be the most critical, while the hf-structure of the $\leftarrow$ 6 line is far too compact to yield any significant shifts."
HESS J18584-020 is a weak gamma-ray source that was reported not to have anv clear cataloged counterpart at any wavelength (Aharonian et al.,HESS J1858+020 is a weak gamma-ray source that was reported not to have any clear cataloged counterpart at any wavelength (Aharonian et al.
 2008)., 2008).
 The nearby radio source G35.6-0.4 was recently re-identified as à SNR (Green 2009)., The nearby radio source G35.6-0.4 was recently re-identified as a SNR (Green 2009).
 ILESS J13534-020 lies towards the southern border of this remnant., HESS J1858+020 lies towards the southern border of this remnant.
 Paron Giacani (2010) have found. using the CO. line from (he Galactic Ring Survey and mid-IB data from the Galactie Legacy Infrared Mid-DPlane survey Extraordinaire (GQLIAIPSE). that there is one or several MCS towards the southern border of SNR G35.6-0.4. likely at the same distance of (he remnant (10.5 kpc).," Paron Giacani (2010) have found, using the $^{13}$ CO line from the Galactic Ring Survey and mid-IR data from the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE), that there is one or several MCs towards the southern border of SNR G35.6-0.4, likely at the same distance of the remnant (10.5 kpc)."
 Paron Giacani (2010) also provide estimates of the clouc’s total molecular mass and density., Paron Giacani (2010) also provide estimates of the cloud's total molecular mass and density.
 They proposed. using a simplified analytical approach described in Torres et al. (," They proposed, using a simplified analytical approach described in Torres et al. ("
2003). Chat hadronic gamuna-ray enmüssion within the clouds. produced by protons diffusing away [rom the SNR. G35.6-0.4. is a possible origin of TESS J1858--020.,"2003), that hadronic gamma-ray emission within the clouds, produced by protons diffusing away from the SNR G35.6-0.4, is a possible origin of HESS J1858+020."
 In a more recent paper. Paron et al. (," In a more recent paper, Paron et al. ("
2011) give more details about the molecular material. obtained via observations with the Atacama Submillimeter Telescope Experiment.,"2011) give more details about the molecular material, obtained via observations with the Atacama Submillimeter Telescope Experiment."
 They discovered a voung stellar object (Y8SO). probably a hieh mass protostar. embedded in the molecular clump. but no evidence of any molecular outflows which mieht in principle reveal the presence of a thermal jet capable of generating the observed eamma-rays.," They discovered a young stellar object (YSO), probably a high mass protostar, embedded in the molecular clump, but no evidence of any molecular outflows which might in principle reveal the presence of a thermal jet capable of generating the observed gamma-rays."
 Paron et al. (, Paron et al. (
2011) concluded again that the most probable origin for the TeV eamuia-rav emission are haclronic interactions between the molecular gas and the cosmic ravs accelerated bv the shock front of SNR G35.6-0.4.,2011) concluded again that the most probable origin for the TeV gamma-ray emission are hadronic interactions between the molecular gas and the cosmic rays accelerated by the shock front of SNR G35.6-0.4.
 Here. we focus on a more in-depth analvsis of this possibility.," Here, we focus on a more in-depth analysis of this possibility."
where a dot denotes a time derivative. IT is the Wnbble constant. σ is the aunihilation cross-section aud vey is the equilibrium value of ay.,"where a dot denotes a time derivative, H is the Hubble constant, $\sigma$ is the annihilation cross-section and $n_{eq}$ is the equilibrium value of $n_{\chi}$."
 Iu the early universe. at high temperature. the last term in this equation dominates and one finds the equiibiuin uuuber deusitv of X particles.," In the early universe, at high temperature, the last term in this equation dominates and one finds the equilibrium number density of $\chi$ particles."
 If lis were always the case then today we would find neslieible numbers of hem aud their energv density would certainly be too little to account for he dark matter., If this were always the case then today we would find negligible numbers of them and their energy density would certainly be too little to account for the dark matter.
" However. as the universe expands it reaches a temperature. shown as thetemperatirc, at which the evolution equation become dominated by the first term on the right- haud side - the damping duc to the he IIubble expansion."," However, as the universe expands it reaches a temperature, known as the, at which the evolution equation become dominated by the first term on the right- hand side - the damping due to the the Hubble expansion."
 After this point. aunibilatious cease aud the distribution of X particles at that time is merely diluted by the expansion at all later times. cading to an abuudance that is uch hieher than the equilibrimm one at those eniperatures.," After this point, annihilations cease and the distribution of $\chi$ particles at that time is merely diluted by the expansion at all later times, leading to an abundance that is much higher than the equilibrium one at those temperatures."
 This is illustrated in figure 1 Ü, This is illustrated in figure \ref{fig:relicabund}~ \cite{Jungman:1995df}.
 Iu fact. to a first approxination. the dark matter abundance reiainiug οαν is given by where Tork Is the typical weak interaction cross-section.," In fact, to a first approximation, the dark matter abundance remaining today is given by where $\sigma_{\rm weak}$ is the typical weak interaction cross-section."
 From this oue cau clearly see why it is that WIMPs ect their naue - weakly interacting particles vield the correct order of iiaguitude to explain the dark matter., From this one can clearly see why it is that WIMPs get their name - weakly interacting particles yield the correct order of magnitude to explain the dark matter.
"We attempted to determine photometric redshifts for the targets listed as ""continuum! and ‘undetected’ in Table 2.. using the photometric redshift code 2002).","We attempted to determine photometric redshifts for the targets listed as `continuum' and `undetected' in Table \ref{spectroscopyjournal}, using the photometric redshift code ."
. In. the cases where continuuntr was detected. we tried to fit a continuum to the observed ÁN. band magnitude. and optical magnitudes derived. from the calibrated: spectra.," In the cases where continuum was detected, we tried to fit a continuum to the observed $K-$ band magnitude, and optical magnitudes derived from the calibrated spectra."
 However. the undetected” sources had no continuum. and therefore an upper limit to the optical bands was caleulatecdl from the spectra ane fi along with the A band magnitude in an attempt to se a minimum redshift at which the optical emission would be below the noise limit.," However, the `undetected' sources had no continuum, and therefore an upper limit to the optical bands was calculated from the spectra and fit along with the $K-$ band magnitude in an attempt to set a minimum redshift at which the optical emission would be below the noise limit."
 We performed the same analysis for sources which hac redshifts. measured. from the spectra., We performed the same analysis for sources which had redshifts measured from the spectra.
 Phe results. show significant discrepancies between the fitted. values fromPEC. and the redshifts measured from spectral features.," The results showed significant discrepancies between the fitted values from, and the redshifts measured from spectral features."
 In some cases. no acceptable fit to the observed colours coulc be found. despite a secure redshift having been determine [rom the spectrum.," In some cases, no acceptable fit to the observed colours could be found, despite a secure redshift having been determined from the spectrum."
 In other cases. theZ-PEG lits were very poorly constrained.," In other cases, the fits were very poorly constrained."
 Furthermore. there was a tendeney. for the fits to congregate around. z2.," Furthermore, there was a tendency for the fits to congregate around $z\sim2$."
 We believe that. photometric redshifts could not. be founcl for several reasons: template mismatch. absence of H1 photometry. and/or dust.," We believe that photometric redshifts could not be found for several reasons: template mismatch, absence of IR photometry, and/or dust."
 First. the active galaxies in our sample mav have a significant contribution from direct or scattered AGN light. especially in the optical bands.," First, the active galaxies in our sample may have a significant contribution from direct or scattered AGN light, especially in the optical bands."
 LW the objects are at 2>1. às expected. from their faint A band magnitudes (see85.2.3). the optical bands trace the rest-frame UV emission. and may also be boosted by voung star formation associated. with the radio jet activity.," If the objects are at $z>1$, as expected from their faint $K-$ band magnitudes (see5.2.3), the optical bands trace the rest-frame UV emission, and may also be boosted by young star formation associated with the radio jet activity."
 In. both cases our galaxies will not be well matched to the template ealaxies inZ-PEGC., In both cases our galaxies will not be well matched to the template galaxies in.
 Second. at the likely redshift range of our sources. L«z4. the Balmer and celiscontinuities shift to the wavelength range ~0.9 to 54m. a range not covered by our spectra.," Second, at the likely redshift range of our sources, $1<z<4$, the Balmer and discontinuities shift to the wavelength range $\sim$ 0.9 to $\mu$ m, a range not covered by our spectra."
 found plausible photometric redshifts in the range 1<2<2 [for seven sources from the 7€ Recshilt Survey using RLJiIN photometry. illustrating the importance of near-infrared photometry for. these objects.," found plausible photometric redshifts in the range $1<z<2$ for seven sources from the 7C Redshift Survey using $RIJHK$ photometry, illustrating the importance of near-infrared photometry for these objects."
 Fhird. it) is cillicult to assess the amount of dust in these galaxies.," Third, it is difficult to assess the amount of dust in these galaxies."
 The photometric redshifts were therefore not. included. in. this paper., The photometric redshifts were therefore not included in this paper.
 The continuum is well detected (Pig. 2)).," The continuum is well detected (Fig. \ref{continua}) ),"
 but we see no emission or absorptions lines., but we see no emission or absorptions lines.
 The rise in the continuum around. 78200 iis probably due to the bbreak at zοI., The rise in the continuum around $\sim$ is probably due to the break at $z\sim 1$.
 Vhe continuum is well detected (Fig. 2)).," The continuum is well detected (Fig. \ref{continua}) ),"
 but we see no emission or absorption lines., but we see no emission or absorption lines.
 The rise in the continuum around. SÜGO00 iis probably due to the break at z~1.2., The rise in the continuum around $\sim$ is probably due to the break at $z\sim 1.2$.
 At 2=3.976. this is the most distant racio galaxy discovered. to date from the SUMSSVSS USS sample.," At $z=3.976$, this is the most distant radio galaxy discovered to date from the SUMSS--NVSS USS sample."
" Phe redshift is based on a single emission ine at Aun,=6051AA.", The redshift is based on a single emission line at $\lambda_{\rm obs}=6051$.
. The continuum cdiscontinuity across the line. and the absence of other lines in our wide spectral coverage identifies his line asLya.," The continuum discontinuity across the line, and the absence of other lines in our wide spectral coverage identifies this line as."
. The redshift is based on a single emission line. which we identify asA3727... based on the absence of confirming lines if the line wereLya.. oorlla.. and the presence of clear underlying continuum emission.," The redshift is based on a single emission line, which we identify as, based on the absence of confirming lines if the line were, or, and the presence of clear underlying continuum emission."
 The redshift is based on a single emission line. which we identify asA3727... based on the absence of confirming lines if the line wereLya.. oorlla.. and the presence of clear underlying continuum emission.," The redshift is based on a single emission line, which we identify as, based on the absence of confirming lines if the line were, or, and the presence of clear underlying continuum emission."
 We observed this source at two different position angles to ensure we covered. all. possible identifications., We observed this source at two different position angles to ensure we covered all possible identifications.
 However. the central object coincident with he radio source remains undetected in our VET spectra.," However, the central object coincident with the radio source remains undetected in our VLT spectra."
 No line or continuum emission was detected in the. 3600s FORS2 spectrum., No line or continuum emission was detected in the 3600s FORS2 spectrum.
 Because of ils extremely stecp radio spectrum. (al=— 1.60) and avourable RA at the time of the FORSL observations. we obtained a further. SIO0S. spectrum.," Because of its extremely steep radio spectrum $\alpha_{843}^{1400}=-1.60$ ) and favourable RA at the time of the FORS1 observations, we obtained a further 8100s spectrum."
 “Phe source remains undetected. indicating either that redshift. determinations of such very faint sources are. bevond the capabilities. of present-day optical spectrographs. or that this source may be at z ZU.," The source remains undetected, indicating either that redshift determinations of such very faint sources are beyond the capabilities of present-day optical spectrographs, or that this source may be at $z\simgt$ 7."
 The redshift is based on a single emission line. which we interpret as Lya.. based. on the," The redshift is based on a single emission line, which we interpret as , based on the"
AR«0.1 (Churchwell2002).. al.1991). (e.g.Welch&Koo2001). (Francoetal.2000:Lizano2008).,"$R< 0.1$ \citep{churchwell02}. \citep{woodchurch89,kurtzetal94}, \citep[e.g.][]{welchetal87,gaumeclaussen90,mehringeretal93,kimkoo01}. \citep{francoetal00,lizano08}."
. ~ (Franco-Ieruáudez," $\sim$ \citep{francheretal04,rodrigetal07,galvmadetal08}."
 of spectral type earlier than B3., of spectral type earlier than B3.
 If the regions were to expand at the sound speed of ionized eas. e;10 is. they would have lifetimes of roughly 10! x3.," If the regions were to expand at the sound speed of ionized gas, $c_i \sim
10$ km/s, they would have lifetimes of roughly $10^4$ yr."
 Less than of an OD stars lifetime of a few μοι. vears should therefore be spent within such a region. so the sale fraction of OB stays should now lie within tle.," Less than of an OB star's lifetime of a few million years should therefore be spent within such a region, so the same fraction of OB stars should now lie within them."
 Tlowever. surveys fud muubers in our Galaxy consistcut with over of OB stars being surrounded by them (Wood&Churclsvell1989:DePreeetal.2005).. or equivaleutlv. lifetimes of ~10? vr if this model is correct.," However, surveys find numbers in our Galaxy consistent with over of OB stars being surrounded by them \citep{woodchurch89,depreeetal05}, or equivalently, lifetimes of $\sim
10^5$ yr if this model is correct."
 A uuuber of explanations have been proposed for this lifetime problem. including confinement iu cloud cores by thermal pressure (DePreectal.1995:Giarcfa-Segura& or turbulent pressure (Xieetal.1996).. rana pressure confinement bv iufall (Yorke19586:Tollenbachetal.1991). or bow shocks (VanBurenetal.1990:MacLowctal.19912:Arthur&Home 2006).. chanrpaegue flows (Ποσοetal.1979:Garcia-Segura&Franco1996:ΑντιToare 2006).. disk evaporation (Wollenbachetal.199D)... and mass-loacect stellar winds (Dysonetal.1995:Redman1996:Williamsotal.1996:Lizanoet 1996).. but most have con argued to have major flaws (MacLowetal.2007).," A number of explanations have been proposed for this lifetime problem, including confinement in cloud cores by thermal pressure \citep{depreeetal95,gsfranco96} or turbulent pressure \citep{xieetal96}, ram pressure confinement by infall \citep{yorke86,hollenbachetal94} or bow shocks \citep{vanburenml90,mlvanburen91,arthurhoare06}, , champagne flows \citep{bodenheimeretal79,gsfranco96,arthurhoare06}, disk evaporation \citep{hollenbachetal94}, and mass-loaded stellar winds \citep{dysonetal95,redmanetal96,williamsetal96,lizanoetal96}, but most have been argued to have major flaws \citep{maclowetal07}."
. We have modeled accretion ou to au ionizing source using three-dimensional sinulatious (Petersetal.2010.jiereafterPaper ID.," We have modeled accretion on to an ionizing source using three-dimensional simulations \citep[][hereafter Paper
I]{petersetal10}."
. These calculations sugeest that accretion can indeed explain the lifetime problem. but iu au unexpected war.," These calculations suggest that accretion can indeed explain the lifetime problem, but in an unexpected way."
 Neto(2002.2007) has argued hat ultracompact and Lypercompact iregious aresimply the ionized portion of an accretion How.," \citet{keto02,keto07} has argued that ultracompact and hypercompact regions aresimply the ionized portion of an accretion flow."
 However. massive stars require accretion af rates excocding LO1 AML. | ," However, massive stars require accretion at rates exceeding $10^{-4}$ $_{\odot}$ $^{-1}$ "
Giant radio galaxies (hereafter GRGs) are luminous but low-surlace brightness sources dominated at radio wavelengths by the emission of extended lobes with linear sizes 1MMpe.,Giant radio galaxies (hereafter GRGs) are luminous but low-surface brightness sources dominated at radio wavelengths by the emission of extended lobes with linear sizes $\gtrsim 1$ Mpc.
" Alorphologically, most of GRGs resemble powerful Fanaroll-Rilev tvpe II objects (FRIs:Fanaroll&Riley1974).. and hence dvnamical models developed for classical doubles are (vpically used to infer (heir intrinsic parameters based on the multilrequency radio data obtained using a variety of radio instruments (e.g..IXonaretal.2008:JamrozyMachalskietal. 2009)."," Morphologically, most of GRGs resemble powerful Fanaroff-Riley type II objects \citep[FR\,IIs;][]{fan74}, and hence dynamical models developed for classical doubles are typically used to infer their intrinsic parameters based on the multifrequency radio data obtained using a variety of radio instruments \citep[e.g.,][]{kon08,jam08,mach09}."
.. The kev diffieultv in modeling and understanding of such sources is however the fact that the main characteristics of the ambient medium surrounding giant lobes are hardly known., The key difficulty in modeling and understanding of such sources is however the fact that the main characteristics of the ambient medium surrounding giant lobes are hardly known.
 In fact. there are almost no observational constraints on the structure. temperature. or density of the intergalactic gas al Alpe distances from (the centers of poor clusters or groups in which powerlul radio galaxies. including GRGs. are expected to be located (Jamrozyetal.2004:Saripalli2005).," In fact, there are almost no observational constraints on the structure, temperature, or density of the intergalactic gas at Mpc distances from the centers of poor clusters or groups in which powerful radio galaxies, including GRGs, are expected to be located \citep{jam04,sar05}."
. But one may also (urn the argument around. and argue thal GRGs can indeed be used of the intergalactic medium (IGM) and of its cosmological evolution at the outskirts. or even outside of groups and clusters of galaxies.," But one may also turn the argument around, and argue that GRGs can indeed be used of the intergalactic medium (IGM) and of its cosmological evolution at the outskirts, or even outside of groups and clusters of galaxies."
 This idea was discussed and pushed forward by. e.g. and Safourisetal.(2009).," This idea was discussed and pushed forward by, e.g., \citet{sub08} and \citet{saf09}."
. Clearly. using GRGs as cosmological probes requires a consensus regarding (he nature of those objects.," Clearly, using GRGs as cosmological probes requires a consensus regarding the nature of those objects."
 Arve Cherefore GRGs just evolved ILE galaxies. for which large linear sizes are strictly due to their advanced ages of the order of LOO MMvr?," Are therefore GRGs just evolved II galaxies, for which large linear sizes are strictly due to their advanced ages of the order of $100$ Myr?"
 Or are they instead intrinsically different [rom classical doubles (assuggestedbv.e.g..INaiser&Alexander1999)?," Or are they instead intrinsically different from classical doubles \citep[as suggested by, e.g.,][]{kai99}?"
? A detailed analvsis of various samples of GRGs [avors the lormer possibility. suggesting however in addition that densities of the medium surrounding the discussed. objects are lower than average lor the whole III population (Mack.etal.1998:SchoenmakersMachalski&Jamrozy 2006).," A detailed analysis of various samples of GRGs favors the former possibility, suggesting however in addition that densities of the medium surrounding the discussed objects are lower than average for the whole II population \citep{mack98,sch00b,mach06}."
. It is interesting to note in this context an analogy with the sources located at the other extremum of the size distribution of radio galaxies. namely Compact Syaunetric Objects (CSOs). which spectroscopically are often classilied as Peaked Spectr (GPS) sources.," It is interesting to note in this context an analogy with the sources located at the other extremum of the size distribution of radio galaxies, namely Compact Symmetric Objects (CSOs), which spectroscopically are often classified as GigaHertz-Peaked Spectrum (GPS) sources."
 For these. the long-debated cuestion was if their linear sizes 0.1]—IO kkpc are due to voung ages (<0.1 MMy) of otherwise regular UI-like radio structures. or instead.are due to exceptionally dense ambient medium frustrating evolved jets and confining their lobes within host galaxies (seeO'Dea1993.lorareview)..," For these, the long-debated question was if their linear sizes $0.1 - 10$ kpc are due to young ages $\leq 0.1$ Myr) of otherwise `regular' II-like radio structures, or insteadare due to exceptionally dense ambient medium frustrating evolved jets and confining their lobes within host galaxies \citep[see][for a review]{odea98}."
 Or clo they constitute vel another intrinsically distinct population of extragalactic radio sources?, Or do they constitute yet another intrinsically distinct population of extragalactic radio sources?
 The merging agreement is that compact radio galaxies are indeed newly born precursors of Classical doubles (andalsooflow-powerFRIsources:seethediscussioninStawarz herein).. even though some of them may reside in particularly overdense environment (see.e.g..Garcia-Burilloetal.2007 )..," The merging agreement is that compact radio galaxies are indeed newly born precursors of classical doubles \citep[and also of low-power FR\,I sources; see the discussion in][and references therein]{sta08}, , even though some of them may reside in particularly overdense environment \citep[see, e.g.,][]{gar07}. ."
The shape of the Initial Mass Function (IME) and its connection to the initial physical conditions in molecular clouds remains one of the fandamental (questions in star formation.,The shape of the Initial Mass Function (IMF) and its connection to the initial physical conditions in molecular clouds remains one of the fundamental questions in star formation.
 More specificallv: Does the voung cluster IME mimic (he integrated field star IME even to very low masses?, More specifically: Does the young cluster IMF mimic the integrated field star IMF even to very low masses?
 Is there a low-mass cut-off to the sub-stellar mass Iunction?, Is there a low-mass cut-off to the sub-stellar mass function?
 The IME over the full range of stellar masses has been extensively studied in the past decades starting with Miller&Sealo(1979).. as well as updates by Ixroupa and Chabrier(2003).," The IMF over the full range of stellar masses has been extensively studied in the past decades starting with \citet{ms79}, as well as updates by \citet{ktg93} and \citet{ch03}."
. It is generally accepted that voung clusters exhibit a IMF above I1 M... flattening out towards low-mass stars.," It is generally accepted that young clusters exhibit a Salpeter-like IMF above 1 $_{\odot}$, flattening out towards low-mass stars."
 Studies of the IAIF have been concentrated on voung clusters. because Chev offer several beneficial characteristics.," Studies of the IMF have been concentrated on young clusters, because they offer several beneficial characteristics."
 Τμοι populations are less likely (ο have undergone significant dynamical mass segregation compared to older clusters. which means (hat a small field can vield a sample representative ol the whole eluster.," Their populations are less likely to have undergone significant dynamical mass segregation compared to older clusters, which means that a small field can yield a sample representative of the whole cluster."
 In addition. voung low mass stars are still located significantly above the main sequence and are (hus several times brighter (han (their main sequence counterparts.," In addition, young low mass stars are still located significantly above the main sequence and are thus several times brighter than their main sequence counterparts."
"(qc 5/3) or Ivaichnan-tvpe(q=3/2). themomentum-dependence becomes a,=1/3 and A,=1/2. respectively,","$q\simeq5/3$ ) or $q=
3/2$ ), themomentum-dependence becomes $\alpha_p = 1/3$ and $\alpha_p=1/2$, respectively."
" Bohin-tvpe diffusion. on the other hand. would imply a,=1. while hard-sphere scattering is described by a,=0."," Bohm-type diffusion, on the other hand, would imply $\alpha_p=1$, while hard-sphere scattering is described by $\alpha_p=0$."
" Note. however. that if one considers electron acceleration by resonant. Langmuir waves. even D,,= const (a,=2) may become possible (Aharonianetal.1986)."," Note, however, that if one considers electron acceleration by resonant Langmuir waves, even $D_p =$ const $\alpha_p=2$ ) may become possible \citep{aharonian86}."
. The svuchrotron energy losses (hat appear in Che second term of Eq. (4)), The synchrotron energy losses that appear in the second term of Eq. \ref{dif}) )
 are In the s-parameter space. (he solution of Eq. (4))," are In the $\gamma$ -parameter space, the solution of Eq. \ref{dif}) )"
" becomes a relativistic Maxwell-like [unction (a,# —1) with and constant. 41 to be defined by the initial conditions.", becomes a relativistic Maxwell-like function $\alpha_p \neq -1$ ) with and constant $A$ to be defined by the initial conditions.
 Note that this is a steady-state solution already including radiative losses and there is no need to invoke extreme values [or the magnetic field., Note that this is a steady-state solution already including radiative losses and there is no need to invoke extreme values for the magnetic field.
" The critical Lorentz [actor , approximately corresponds to (he energy ab which acceleration on timescale is balanced by (synchrotron) cooling on timescale (ou.=1/|5.p].", The critical Lorentz factor $\gamma_c$ approximately corresponds to the energy at which acceleration on timescale is balanced by (synchrotron) cooling on timescale $t_{\rm cool}=1/[\beta_s p]$.
 Depending on the choice of parameters. a relatively large range of values for 5. is possible aud thus. cut-off energies of the order of 2.cLO? may well be achieved.," Depending on the choice of parameters, a relatively large range of values for $\gamma_{c}$ is possible and thus, cut-off energies of the order of $\gamma_{c}\sim 10^{5}$ may well be achieved."
 Consider. [or example. Bohin-ivpe diffusion with. 7=yry/e. ry=5mc>(eB)B the electron gvro-radius. and à>1.," Consider, for example, Bohm-type diffusion with $\tau =\eta r_g/c$, $r_g=\gamma m_e c^2/(e B)$ the electron gyro-radius and $\eta \geq 1$."
" Using+ lace=leoot the maximun electron Lorentz [actor becomes 5,2105(04/0.010)ο”..."," Using $t_{\rm acc} =t_{\rm cool}$, the maximum electron Lorentz factor becomes $\gamma_c \simeq
10^6~(v_A / 0.01 c)~(1~\mathrm{G}/B)^{1/2}\eta^{-1/2}$."
 The svnchrotron spectrum (hat arises from a Maxwell-like electrondistribution is dominated bv the emission of electrons wilh 5. (Fie. 3))., The synchrotron spectrum that arises from a Maxwell-like electrondistribution is dominated by the emission of electrons with $\gamma_{c}$ (Fig. \ref{SSCmax}) ).
" It exhibits the characteristic 1/3-slope up to (he corresponding. oe“svuchrotron cut-offB frequency”BEM fis""""~obsz-cay where b=D/D,, and bsocneeeh."," It exhibits the characteristic $1/3$ -slope up to the corresponding ""synchrotron cut-off frequency"" $h \nu^{\rm syn}_{c}\sim\delta b\gamma^{2}_{c}$ where $b=B/B_{cr}$ and $B_{cr}=m^2c^3/e\hbar$."
 Thus the Compton spectrum is very similar to the one resulting. [rom a narrow power-law if one chooses a value [or the cut-off enerev close to (he minimum electron energv of the power-law distribution., Thus the Compton spectrum is very similar to the one resulting from a narrow power-law if one chooses a value for the cut-off energy close to the minimum electron energy of the power-law distribution.
 The peak of the Compton flux then contains information for the cut-off energy as PraX Je , The peak of the Compton flux then contains information for the cut-off energy as $\nu^{c}_{\rm peak}\propto \gamma_{c}$ .
Note that for an electron distribution of the form of eq. (6)), Note that for an electron distribution of the form of eq. \ref{max}) )
 that exhibits an exponential eutolf xexp[7(5/5.) 7). the corresponding cut-olf in the synchrotron spectrum appears much," that exhibits an exponential cutoff $\propto
\exp[-(\gamma/\gamma_{c})^{\beta}]$ , the corresponding cut-off in the synchrotron spectrum appears much"
Svensson (1996). in NSPEC. to calculate the spectrum. produced by thermal Comptonization of photons emitted. with a disk blackbody spectrum ane scattered. in a hot corona.,"Svensson (1996), in XSPEC, to calculate the spectrum produced by thermal Comptonization of photons emitted with a disk blackbody spectrum and scattered in a hot corona."
 The model parameters allowed. to vary. in the fits are the disk temperature Adin. the corona electron temperature Ad). and the corona optical depth 7.," The model parameters allowed to vary in the fits are the disk temperature $kT_{\rm in}$, the corona electron temperature $kT_{e}$, and the corona optical depth $\tau$."
 Because the metallicity of Holmboerg LL is significantly lower than solar. we use two absorption components: one to model absorption within the Milky Way for which we fix the metallicity at solar and fix the absorption column density Ny=(342£0.53)107emi and a second to model absorption within Lolmbere LL for which wefix the metallicity at Z=0.07Z. (Mirioni2002).," Because the metallicity of Holmberg II is significantly lower than solar, we use two absorption components: one to model absorption within the Milky Way for which we fix the metallicity at solar and fix the absorption column density $N_{\rm H} = (3.42 \pm 0.3) \times 10^{20}
\rm \, cm^{-2}$, and a second to model absorption within Holmberg II for which wefix the metallicity at $Z = 0.07 Z_{\sun}$ \cite{mirioni02}."
. We performed a simultaneous fit to the data for all three observations in which the absorption within Llolmberg LE was the same for all observations and the Comptonization model paramoeters were allowed to vary incliviclually for cach observation., We performed a simultaneous fit to the data for all three observations in which the absorption within Holmberg II was the same for all observations and the Comptonization model parameters were allowed to vary individually for each observation.
 We [ound an adequate fit with yo /Dok = 273.5/285., We found an adequate fit with $\chi^2$ /DoF = 273.5/285.
 The best fit parameters are reported in Table 1.., The best fit parameters are reported in Table \ref{specfits}.
 Phe best fit column density is (3.720.5)107+em.27.," The best fit column density is $(3.7 \pm 0.5) \times 10^{21} \rm \, cm^{-2}$."
 We note that this column density is significantly above that found by Dewangan et ((2004) because the absorption. bevond. the Galactic component. is caleulated for the low metallicity appropriate o Lolmbere LL.," We note that this column density is significantly above that found by Dewangan et (2004) because the absorption, beyond the Galactic component, is calculated for the low metallicity appropriate to Holmberg II."
 The three XNMM-Newton observations cover the extremes of N-rav. Lux detected. from. Holmberg Il N-1 (Dewanganetal.2004)., The three XMM-Newton observations cover the extremes of X-ray flux detected from Holmberg II X-1 \cite{dewangan04}.
.. Pherefore. these observations also ikelv cover the range of X-ray spectral variations in the ποσο," Therefore, these observations also likely cover the range of X-ray spectral variations in the source."
 We note a thermal bremsstrahlung model as used. by Pakull Mirioni (2002) for input to their photoionization modelling does not provide an adequate Π to any of he observations. even with use of two clistinet absorption COMPnents.," We note a thermal bremsstrahlung model as used by Pakull Mirioni (2002) for input to their photoionization modelling does not provide an adequate fit to any of the observations, even with use of two distinct absorption components."
 Observations were mace of Holmberg LL centered. on the position of the ULX using the Acwancec Camera for Surveys (ACS) on the Hubble Space Telescope (LIST) under CO program 9684 (PL WKaaret)., Observations were made of Holmberg II centered on the position of the ULX using the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST) under GO program 9684 (PI Kaaret).
" Observations weremace in the narrow band. filters ΤΗΝ centered. onHenr A4686. FRS505N centered. on A4861. and E14656N. centered on Ol, A6300."," Observations weremade in the narrow band filters FR462N centered on $\lambda$ 4686, FR505N centered on $\beta$ $\lambda$ 4861, and FR656N centered on ] $\lambda$ 6300."
. Images in the FR4AG2N filter. were obtained on 21 January 2003 and 25 June 2003., Images in the FR462N filter were obtained on 21 January 2003 and 25 June 2003.
 For the January observation. the aimpoint was placed closed to the ramp filter οσο and there ave non-uniformities in. transmission across the image.," For the January observation, the aimpoint was placed closed to the ramp filter edge and there are non-uniformities in transmission across the image."
 For this reason. we quote Duxes for the emission based onlv on the June observation.," For this reason, we quote fluxes for the emission based only on the June observation."
 The images in all of the other filters were obtained on 24 November 2002., The images in all of the other filters were obtained on 24 November 2002.
 All of the narrow band filters have a bandwidth., All of the narrow band filters have a bandwidth.
 The recession velocity of Llolmabere LL is 157 km/s (Straussetal.1992).. so the redshifted emission lines lie within the filter bandpasses.," The recession velocity of Holmberg II is 157 km/s \cite{strauss92}, so the redshifted emission lines lie within the filter bandpasses."
 La addition. observations were made in the medium filters FIU59M and 1550M for continuum imaging and continuum subtraction.," In addition, observations were made in the medium filters FR459M and F550M for continuum imaging and continuum subtraction."
 The standard processing for ACS data does not perform cosmic-ray removal for images without cosmic-ray splits., The standard processing for ACS data does not perform cosmic-ray removal for images without cosmic-ray splits.
 Beeause all of our observations except those in the ΕΠΗΝ filter were performed in dither patterns without cosmic-ray splits. we re-processecl all observations using the4 task in SVSDAS 3.1 to remove the cosmic rav hits.," Because all of our observations except those in the FR462N filter were performed in dither patterns without cosmic-ray splits, we re-processed all observations using the task in STSDAS 3.1 to remove the cosmic ray hits."
" We found. residual sky level olfsets in the images ancl removed these by fitting a gaussian. to those pixels not containing astronomical objects in a 30"". field centered near the ULNA and. subtracting olf the gaussian centroid.", We found residual sky level offsets in the images and removed these by fitting a gaussian to those pixels not containing astronomical objects in a $30\arcsec \times 30\arcsec$ field centered near the ULX and subtracting off the gaussian centroid.
 We aligned the F550M. (narrow V. band) image o 10 stars selected. from the ΝΟ A2.0 catalog (Monetetal.1996). using the tool from the Smithsonian Astrophysical Observatory Telescope Data Center., We aligned the F550M (narrow V band) image to 10 stars selected from the USNO A2.0 catalog \cite{monet96} using the tool from the Smithsonian Astrophysical Observatory Telescope Data Center.
 Dased on he residual olfsets for the LO stars. we estimate that the astrometric uncertainty is 0.37.," Based on the residual offsets for the 10 stars, we estimate that the astrometric uncertainty is $0.3\arcsec$."
 We then aligned each other image to the aspect corrected E550M image using theAL ools ancl Clodyctal.1993)., We then aligned each other image to the aspect corrected F550M image using the tools and \cite{tody93}.
.. We checked he alignment using the toolaregisfer and found that the residual misalignments were less than 0.1 pixel., We checked the alignment using the tool and found that the residual misalignments were less than 0.1 pixel.
 We produced. continuum subtracted images using the 1I459M image to estimate the continuum for the ΕΠΟΝ A4686 and E15505N LL2/.A4861 images. and the E550D»a image to estimate the continuum for the FR6G56N O1] A6300 image.," We produced continuum subtracted images using the FR459M image to estimate the continuum for the FR462N $\lambda$ 4686 and FR505N $\beta$ $\lambda$ 4861 images, and the F550M image to estimate the continuum for the FR656N ] $\lambda$ 6300 image."
 Since we are interested primarily in the cilfuse. nebular emission. we located. the stars in cach frame ancl subtracted olf the stellar emission. before. performing the continuum subtraction.," Since we are interested primarily in the diffuse, nebular emission, we located the stars in each frame and subtracted off the stellar emission before performing the continuum subtraction."
 We fit a Mollat. profile to several bright stars to determine the point spread. function shape. and then used that fixed shape in fitting and subtracting the stars.," We fit a Moffat profile to several bright stars to determine the point spread function shape, and then used that fixed shape in fitting and subtracting the stars."
 We note that the ΤΠΕΟΝ image used for continuum subtraction ofthe FRAG2ZN image contains the line., We note that the FR459M image used for continuum subtraction of the FR462N image contains the line.
 We correct for the apparent reduction in the line Dux caused by partial subtraction of the line as described below., We correct for the apparent reduction in the line flux caused by partial subtraction of the line as described below.
" For the continuum subtraction of the nebula. we assumed: that. the intrinsic. continuum. spectrum: is. flat. BAYxA"" with »=0 and reddened with an extinction of E(B-V) = 0.024 (Stewartetal. 2000).."," For the continuum subtraction of the nebula, we assumed that the intrinsic continuum spectrum is flat, $F(\lambda) \propto \lambda^{n}$ with $n = 0$ and reddened with an extinction of E(B-V) = 0.024 \cite{stewart00}. ."
 For the and 1.7 images. the continuum band lies close to the line wavelength. and changing the continuum slope does not. significantly," For the and $\beta$ images, the continuum band lies close to the line wavelength, and changing the continuum slope does not significantly"
"and in the others there are estimates of only one: for these we assume that 0,=6).","and in the others there are estimates of only one: for these we assume that $\theta_t =
\theta_j$."
 This is reasonable on the basis of the similarity where there are independent estimates. ancl by the fact that many (though not all) images show that jets are roughly aligned with the axes of the tori when seen in projection.," This is reasonable on the basis of the similarity where there are independent estimates, and by the fact that many (though not all) images show that jets are roughly aligned with the axes of the tori when seen in projection."
 The uncertainties adopted for the inclination angles given in Table 2 are of two (vpes., The uncertainties adopted for the inclination angles given in Table 2 are of two types.
 In specific cases noted in the footnotes to the table. the values of A@ are taken from the relerences cited.," In specific cases noted in the footnotes to the table, the values of $\Delta\theta$ are taken from the references cited."
 For the others we have mace rough estimates. based on the the observations and. where relevant. ou the agreement between different estimates: for simplicity in (hese cases we have adopted values for A@ of 75° or z10.," For the others we have made rough estimates, based on the the observations and, where relevant, on the agreement between different estimates: for simplicity in these cases we have adopted values for $\Delta\theta$ of $\pm5\arcdeg$ or $\pm10\arcdeg$."
" For the sources where we assume 0,=@,. (he errors in Fig."," For the sources where we assume $\theta_t = \theta_j$, the errors in Fig."
 1 are correlated: and for the three objects in which /; is estimated directly [rom proper motions. (he results are independent of the inclination.," 1 are correlated; and for the three objects in which $t_j$ is estimated directly from proper motions, the results are independent of the inclination."
 KjPu 8., KjPn 8.–
" 0, is re-determined from the velocity strip maps of the torus 1998).. and agrees with 0; from the jet kinematics given by Meaburn(1997)."," $\theta_t$ is re-determined from the velocity strip maps of the torus \citep{for98}, , and agrees with $\theta_j$ from the jet kinematics given by \citet{mea97}."
. /; is from optical proper motions of the jets. 242:3 mas ! (Meaburn L997)..," $t_j$ is from optical proper motions of the jets, $34\pm3$ mas $^{-1}$ \citep{mea97}. ."
 AL 1-10., M 1-16.–
 The adopted value of 0; (romSchwarz1992). is at the top of the range given bv Corradi&Schwarz(1993): this value is prelerred. because the jets align with the torus axis on the skv ancl this is (he closest value to 8;. which is well determined.," The adopted value of $\theta_j$ \citep[from][]{sch92} is at the top of the range given by \cite{cor93}; this value is preferred because the jets align with the torus axis on the sky and this is the closest value to $\theta_t$, which is well determined."
 The inner edge of molecular torus is not well resolved by the CO observations: the limit in Fig., The inner edge of molecular torus is not well resolved by the CO observations; the limit in Fig.
 2 corresponds to the radius of the small ionized nebula., 2 corresponds to the radius of the small ionized nebula.
 /; for the (wo additional jet components in Fig., $t_j$ for the two additional jet components in Fig.
 2 are 1050 vr and 740 vr (Schwarz1992)., 2 are 1050 yr and 740 yr \citep{sch92}.
. M 2-9., M 2-9.–
 4; is from optical proper motions of the jets. 51-7 mas vr.! )..," $t_j$ is from optical proper motions of the jets, $51\pm7$ mas $^{-1}$ \citep{sch97}."
 AL 1-92., M 1-92.–
 The adopted value of 9;=51325 (S0lEL994) is from a geometrical method using Doppler shifts in the jets.," The adopted value of $\theta_j=57\pm5\arcdeg$ \citep{sol94}
 is from a geometrical method using Doppler shifts in the jets."
" /; and /; are based on the equatorial and polar velocity gradients (12 and 7.6 +. respectively) given by Alcoleaetal.(2007): using the gradient for the torus gives a value of /; that depends on (an2,."," $t_t$ and $t_j$ are based on the equatorial and polar velocity gradients (12 and 7.6 $^{-1}$, respectively) given by \cite{alc07}; using the gradient for the torus gives a value of $t_t$ that depends on $\tan
\theta_t$."
 The gradient suggests a torus ejection event. although the jet oulllows in this source are particularly wide-angle and extend to low latitudes so that the torus may be winc-swept: this may account for the evadient.," The gradient suggests a torus ejection event, although the jet outflows in this source are particularly wide-angle and extend to low latitudes so that the torus may be wind-swept; this may account for the gradient."
 Aleoleaetal.(2007) assume the torus and jets are ejected at the same time. which requires an inclination angle of 51.5° to get the same expansion (me scales.," \cite{alc07} assume the torus and jets are ejected at the same time, which requires an inclination angle of $51.5\arcdeg$ to get the same expansion time scales."
 In Figs.1 and 2we use the independent value of Solf(1994). eiven above., In Figs.1 and 2we use the independent value of \cite{sol94} given above.
centered al z22withastandarddeviationof1.,centered at $=$ 2 with a standard deviation of 1.
M9 hilethisisnolexacllylhesameasthedistribulionsseenincosim ," While this is not the same as the distributions seen in cosmological models, it is similar enough, and the exact shape does not effect the conclusion."
"by simply applving the redshift correction: Ej,=Epearons(l+2)."," With these two values, we first find the intrinsic $E_{peak}$ by simply applying the redshift correction: $E_{peak}=E_{peak,obs}(1+z)$."
 We then use the Amati relation to derive {οτο. and use equation 2 to get the observed ρω.," We then use the Amati relation to derive $E_{\gamma, iso}$, and use equation 2 to get the observed $S_{bolo}$."
 As such. the figure shows a realistic distribution. or at least [or no measurement uncertainties.," As such, the figure shows a realistic distribution, or at least for no measurement uncertainties."
 In the figure. we see (hat there are no violators (1.e.. bursts appearing below the Amati limit). with most bursts appearing close to the limit line.," In the figure, we see that there are no violators (i.e., bursts appearing below the Amati limit), with most bursts appearing close to the limit line."
 This figure is a central illustration of the Nakar Pian test. which we will extend in (his paper.," This figure is a central illustration of the Nakar Piran test, which we will extend in this paper."
" If we allow for ordinary scatter caused by measurement errors in. £4, and 5j4,. then the tight scatter in Figure 2 is lost."," If we allow for ordinary scatter caused by measurement errors in $E_{peak,obs}$ and $S_{bolo}$ , then the tight scatter in Figure \ref{fig:NaPAmatiMC} is lost."
 This is shown in Figure 3.. where suddenly somewhat less than hall of the bursts become violators.," This is shown in Figure \ref{fig:NaPAmatiMCScatter}, where suddenly somewhat less than half of the bursts become violators."
 For (his simulation. we assumed that. the measurement errors have a log-normal distribution with a one-sigma width of 0.25 (Collazzi et al.," For this simulation, we assumed that the measurement errors have a log-normal distribution with a one-sigma width of 0.25 (Collazzi et al."
 2011)., 2011).
 The exact fraction of violators will depend on the size of the observational scatter., The exact fraction of violators will depend on the size of the observational scatter.
 In this realistic simulation. ~40% of the bursts are below the Amati limit line.," In this realistic simulation, $\sim$ of the bursts are below the Amati limit line."
 The point of this figure is (hat normal and expected observational measurement errors will lead to nearly half the bursts being apparent violators., The point of this figure is that normal and expected observational measurement errors will lead to nearly half the bursts being apparent violators.
 Importantlyv. (his scatter does not explain the hieh violator rates reported by Band Preece (2005) and Goldstein et al. (," Importantly, this scatter does not explain the high violator rates reported by Band Preece (2005) and Goldstein et al. ("
2010).,2010).
 This discrepancy is the main topic of this paper., This discrepancy is the main topic of this paper.
" For comparison. we can also consider how the Shoe—E,4,5;, diagram would look if neither the Amati or Ghirlanda relations were valid."," For comparison, we can also consider how the $S_{bolo} - E_{peak,obs}$ diagram would look if neither the Amati or Ghirlanda relations were valid."
 For this. we have constructed another Alonte Carlo simulation (see Figure 4)).," For this, we have constructed another Monte Carlo simulation (see Figure \ref{fig:NaPDistMC}) )."
 As in Figure 2. we have assumed no measurement errors. no selection bv satellite detectors. and we have adopted realistic laminosity and distauce distributions. bul we have mace no constraints from either the Amati or Ghirlanda relations.," As in Figure 2, we have assumed no measurement errors, no selection by satellite detectors, and we have adopted realistic luminosity and distance distributions, but we have made no constraints from either the Amati or Ghirlanda relations."
 We start by selecting burst distances and energies in the 100-500 keV such that they reproduce the observed Ιουν)—Ιου) curves for.DATSE (Fenimore et al., We start by selecting burst distances and energies in the 100-500 keV such that they reproduce the observed $\log(N)-\log(P)$ curves for (Fenimore et al.
 1993. Fishinan Meegan 1995).," 1993, Fishman Meegan 1995)."
 We then generate [τρως based on a log normal distribution with some loose connection to the brightness of the burst (as seen in Mallozzi et al., We then generate $E_{peak}$ based on a log normal distribution with some loose connection to the brightness of the burst (as seen in Mallozzi et al.
 1995)., 1995).
 We then apply a bolometric correction with (a=—1.0 and ~]= —2.0)., We then apply a bolometric correction with $\alpha=-1.0$ and $\beta=-2.0$ ).
 The result is an in illustration of the intrinsic distribution of bursts on the sky., The result is an in illustration of the intrinsic distribution of bursts on the sky.
 Our simulation of 10.000 bursts has approximate edges at 20 and 3000 keV. plus lower and upper edees simply where we cul olf the logCN)—log(P?) curve.," Our simulation of 10,000 bursts has approximate edges at 20 and 3000 keV, plus lower and upper edges simply where we cut off the $\log(N)-\log(P)$ curve."
" The kev point is that Figures 2 and + are greatly. different. because low-fIuence bursts will dominate unless some law/correlation forces these low-fluence events to have low-£E,,,5,5."," The key point is that Figures 2 and 4 are greatly different, because low-fluence bursts will dominate unless some law/correlation forces these low-fluence events to have $E_{peak,obs}$."
" 90 we have (wo extreme cases that produce greatly different distributions in the Spor—E, diagram."," So we have two extreme cases that produce greatly different distributions in the $S_{bolo} - E_{peak,obs}$ diagram."
 Both Figures 2 and 4 are For (he intrinsic distributions of GRBs in a realistic case wilh no effects of detector thresholds or measurement uncertainties., Both Figures 2 and 4 are for the intrinsic distributions of GRBs in a realistic case with no effects of detector thresholds or measurement uncertainties.
 From a comparison of Figures 2 and 3. we see that realistic measurement errors will substantially smearthe underlving," From a comparison of Figures 2 and 3, we see that realistic measurement errors will substantially smearthe underlying"
"At z>1.65, the Lyman-o line shifts above the atmosphericcutoff at ~3000A, and ground based surveys can efficiently search for DLA systems.","At $z>1.65$, the $\alpha$ line shifts above the atmosphericcutoff at $\sim3000\ang$, and ground based surveys can efficiently search for DLA systems."
 Several such surveys have determined the z>2 neutral gas density (Wolfeetal.1995;Storrie-Lombardi&terdaemeetal.," Several such surveys have determined the $z>2$ neutral gas density \citep{ Wol95, SL00, Pro05, Not09b}."
" At lower redshifts, space-based spectra are necessary 2009)..to measure the neutral hydrogen column densities in the UV."," At lower redshifts, space-based spectra are necessary to measure the neutral hydrogen column densities in the UV."
" As DLAs (and QSOs that are bright enough in the far ultraviolet (FUV) to be accessible to previous UV are relatively rare, the number of such systems spectrographs)currently known at z« lis small compared to available samples at high redshift."," As DLAs (and QSOs that are bright enough in the far ultraviolet (FUV) to be accessible to previous UV spectrographs) are relatively rare, the number of such systems currently known at $z<1$ is small compared to available samples at high redshift."
 The Cosmic Origins Spectrograph is à new instrument package (Froning&Green(COS)2009) installed on the Hubble Space Telescope that enables us to study these absorbers at low redshift with unprecedented efficiency.," The Cosmic Origins Spectrograph (COS) is a new instrument package \citep{Fro09}
 installed on the Hubble Space Telescope that enables us to study these absorbers at low redshift with unprecedented efficiency."
" As z«0.5 spans ~ 40 percent of the age of the universe, the low redshift absorbers accessible with COS are crucial for understanding cosmic chemical evolution and the cosmological gas mass density and for linking their properties to their higher redshift counterparts."," As $z<0.5$ spans $\sim$ 40 percent of the age of the universe, the low redshift absorbers accessible with COS are crucial for understanding cosmic chemical evolution and the cosmological gas mass density and for linking their properties to their higher redshift counterparts."
" In this paper, we report on the first DLA systems observed with COS."," In this paper, we report on the first DLA systems observed with COS."
" We present a detailed analysis of one such system, a DLA in the line of sight to the QSO SDSS J10094-0713, to illustrate the scientific potential of the observations."," We present a detailed analysis of one such system, a DLA in the line of sight to the QSO SDSS J1009+0713, to illustrate the scientific potential of the observations."
 The structure of the paper is as follows., The structure of the paper is as follows.
 In $2 we describe in general the HST program and data reduction methods we have used for the COS spectra., In $\S$ 2 we describe in general the HST program and data reduction methods we have used for the COS spectra.
" In 8 3 we describe the observations of the field of SDSS J1009+0713, including the ground based spectra of the QSO, the COS UV spectra, Keck/HIRESand imaging of the field with the Wide Field Camera 3 (WFC3)."," In $\S$ 3 we describe the observations of the field of SDSS J1009+0713, including the ground based Keck/HIRES spectra of the QSO, the COS UV spectra, and imaging of the field with the Wide Field Camera 3 (WFC3)."
" In § 4 we derive chemical abundances for this system, discuss the physical state of the gas as determined by the CI and C II* lines, and discuss the properties of the galaxies in the field as seen in the WFC3 images."," In $\S$ 4 we derive chemical abundances for this system, discuss the physical state of the gas as determined by the CI and C II* lines, and discuss the properties of the galaxies in the field as seen in the WFC3 images."
" In § 5 we combine the measurements of N(H from other absorbers in this program to measure the I)cosmological mass density of neutral gas, Og1, at z«0.35 in a blind survey with Az~12."," In $\S$ 5 we combine the measurements of N(H I) from other absorbers in this program to measure the cosmological mass density of neutral gas, $\Omega_{\rm H \ I}$, at $z<0.35$ in a blind survey with $\Delta z\sim12$."
 Conclusions are summarized in 8 6., Conclusions are summarized in $\S$ 6.
" Throughout this paper, we adopt a cosmological model with Qm=0.30, Q4—0.70, and Ηρ--τθ km s! Mpc-!."," Throughout this paper, we adopt a cosmological model with $\Omega_m=0.30$, $\Omega_{\Lambda}=0.70$, and $_0$ =70 km $^{-1}$ $^{-1}$ ."
" The targets presented here were observed as part of program GO11598, a program focused on studying multiphase baryons in the halos of L=L* galaxies at 15-150 kpc impact parameters."," The targets presented here were observed as part of program GO11598, a program focused on studying multiphase baryons in the halos of $L\ga L^{\star}$ galaxies at 15-150 kpc impact parameters."
" Target QSOs for this program were selected based on sufficient FUV flux and the presence of a galaxy seen in the Sloan Digital Sky Survey (SDSS, Yorketal. 2000)) with impact parameter <150 kpc and a spectroscopic or photometric redshift p0.15«z 0.35."," Target QSOs for this program were selected based on sufficient FUV flux and the presence of a galaxy seen in the Sloan Digital Sky Survey (SDSS, \citealt{York00}) ) with impact parameter $\rho<150$ kpc and a spectroscopic or photometric redshift $0.15<z<0.35$ ."
 Sightlines with strong Mg II systems seen in the SDSS spectra were avoided due to the possibility ofa Lyman-limit system (LLS) at z <0.35, Sightlines with strong Mg II systems seen in the SDSS spectra were avoided due to the possibility ofa Lyman-limit system (LLS) at $z\la$ 0.35
Iu equation 3.. Aer.y) is the telescope pupil fiction (a mask showing the obscuration) aud the phase difference εαν) has the dimension of leneth.,"In equation \ref{eqn_p}, , $A(x,y)$ is the telescope pupil function (a mask showing the obscuration) and the phase difference $\psi(x,y)$ has the dimension of length."
 The pixel scale of the PSF array eiven by equation | ds simply FA/2. where Fis the fratio.," The pixel scale of the PSF array given by equation \ref{eqn_psf} is simply $F \lambda/2$, where $F$ is the f-ratio."
 Figure 5. displavs the examples of the LSST PSFs geucrated in this way at differcut integration times., Figure \ref{fig_psf_time_evolution} displays the examples of the LSST PSFs generated in this way at different integration times.
 The PSF at £=0 shows the typical instantaneous speckle., The PSF at $t=0$ shows the typical instantaneous speckle.
 In this figure we do not include either the optical aberration or the charge diffusion by CCDs. the effects of which we however later add to generate the simulated LSST nuages.," In this figure we do not include either the optical aberration or the charge diffusion by CCDs, the effects of which we however later add to generate the simulated LSST images."
 As the exposure time inereases. more speckles are stacked together. which makes the resulting PSF rouuder aud the ireeular features preseut in the individual speckles more smeared (cle Vries ct al.," As the exposure time increases, more speckles are stacked together, which makes the resulting PSF rounder and the irregular features present in the individual speckles more smeared (de Vries et al."
 2007)., 2007).
 A quantitative study on the impact of the atmospheric turbulence on the ellipticity aud its spatial correlation is needed to support the validitv of our simulation hereafter., A quantitative study on the impact of the atmospheric turbulence on the ellipticity and its spatial correlation is needed to support the validity of our simulation hereafter.
" As is discussed im roefscctiongocellane, icinodel LSST PSF eariationC CCDwithpolgiomials."," As is discussed in \\ref{section_focal_plane}, we model LSST PSF variation CCD-by-CCD with polynomials."
"I ftheauisotropiepowerf roithcatsyoeyfieationtCD ος docabqrlanct corpo fealwcfttoonf ibd dela gesoi smooththeinherent PSF variation,"," If the anisotropic power from the atmosphere within 15 sec exposure turns out to be too strong and changes on a very small scale, any conventional interpolation scheme with the finite number of stars will over-smooth the inherent PSF variation."
 We define the ellipticity of PSE as(«.6)/(a|b). where a and b ave the semi-major and auiuor axes of the PSFs. respectively.," We define the ellipticity of PSF as $(a-b)/(a+b)$, where $a$ and $b$ are the semi-major and -minor axes of the PSFs, respectively."
 In general. the PSF isophotes change with radius.," In general, the PSF isophotes change with radius."
 Thus. the measurement somewhat depeuds ou the weighting scheme.," Thus, the measurement somewhat depends on the weighting scheme."
" We measure the ellipticity of the PSF using the following quadrupole moments: where £(@) is the pixel intensity at @. 0;j, is the couter of the star. aud (8) is the optimal weight function required to suppress the noise in the outskirts."," We measure the ellipticity of the PSF using the following quadrupole moments: where $I(\symvec{\theta})$ is the pixel intensity at $\symvec{\theta}$, $\bar{\theta}_{i(j)}$ is the center of the star, and $W(\symvec{\theta})$ is the optimal weight function required to suppress the noise in the outskirts."
" For the dciffraction-Inuited PSF of LSST (Figure 8)). we choose a Gaussian with a FWIIM of —0.3""7foy (8) whereas for the turbuleuce-Imited PSF the EWIIM of the Caussiau weight ΠιοΊο used is ~0.7"","," For the diffraction-limited PSF of LSST (Figure \ref{fig_ellipticity_focal_plane_no_atmos}) ), we choose a Gaussian with a FWHM of $\sim$ for $W(\symvec{\theta})$ whereas for the turbulence-limited PSF the FWHM of the Gaussian weight function used is $\sim0.7\arcsec$."
 The quadrupole moments are converted to the ellipticitv € via: The magnitude of the stick is proportional to εἰ=V6εἰ. and the orientation auele is eiven bw 0.5tanle€|).," The quadrupole moments are converted to the ellipticity $\symvec{\epsilon}$ via: The magnitude of the stick is proportional to $|\symvec{\epsilon}|=\sqrt{\epsilon_+^2 + \epsilon_{\times}^2}$ , and the orientation angle is given by $0.5\tan^{-1}( \epsilon_{\times}/ \epsilon_{+})$."
 Figure 6aa displavs the PSF ellipticity variation for the 15 s exposure within a Ik«1k CCD of LSST bx atmosphere., Figure \ref{fig_at_psf}a a displays the PSF ellipticity variation for the 15 s exposure within a $\times$ 4k CCD of LSST by atmosphere.
 Because no optical aberration is imtroduced vet. this £000. “whiskers” show the purely atimospheric contribution.," Because no optical aberration is introduced yet, this 4000 “whiskers” show the purely atmospheric contribution."
 The size of the sticks represents the magnitude of the ellipticity whereas the the oricutation of the stick is aligned with the position anele of the elongation., The size of the sticks represents the magnitude of the ellipticity whereas the the orientation of the stick is aligned with the position angle of the elongation.
 A clear spatial correlation is visible on a scale of —1 (~300 pixels)., A clear spatial correlation is visible on a scale of $\sim$ $\arcmin$ $\sim$ 300 pixels).
 However. the average magnitude is less than .. aud thus the resulting anisotropic power is not strong.," However, the average magnitude is less than , and thus the resulting anisotropic power is not strong."
" Figure 6bb shows the cllipticity correlation based on the following equation: where the product e(7)ο|r) d performed for cach pair separated by r. and the e; aud ο. are the tanecutial aud the 15? components.respectively, witli respect to the line connecting the pair."," Figure \ref{fig_at_psf}b b shows the ellipticity correlation based on the following equation: where the product $e (r^{\prime} )~ e (r^{\prime} + r)$ is performed for each pair separated by $r$, and the $e_{t}$ and $e_{\times}$ are the tangential and the $\degr$ components,respectively, with respect to the line connecting the pair."
 The amplitude of the correlation function decreases with exposure time., The amplitude of the correlation function decreases with exposure time.
 From the comparison with the contribution due to the telescope aud camera aberration (see retsectiongocallune)). weconeludethatthesmallscalea isotropic pou," From the comparison with the contribution due to the telescope and camera aberration (see \\ref{section_focal_plane}) ), we conclude that the small scale anisotropic power for the delivered PSF is not dominated by the atmospheric turbulence."
 The 61 c, These results have been validated with a series of 15 sec exposures of the globular cluster NGC2419 on the Subaru telescope.
m diameter focal plane of LSST will be tiled with 189 Ik< lk CCDs., The 64 cm diameter focal plane of LSST will be tiled with 189 $\times$ 4k CCDs.
 As-built heights might vary up to ~10 inicrous (peak-to-peak) iu a complicated wav relative to the nominal flat surface., As-built heights might vary up to $\sim10$ microns (peak-to-peak) in a complicated way relative to the nominal flat surface.
 Although this flatuess deviation is small compared with other canieras dn existing facilities. the αμα. fratio of the iustruinent makes this focal plane flatuess variation plav a critical role in aberration-induced PSF behavior (deptl of focus X £ratio).," Although this flatness deviation is small compared with other cameras in existing facilities, the small f-ratio of the instrument makes this focal plane flatness variation play a critical role in aberration-induced PSF behavior (depth of focus $\propto$ f-ratio)."
 In Table 1 we stuumarize the current issued. to vendors. as well as the values used for the current simulation.," In Table 1 we summarize the current specification of the focal plane error budget originally issued to vendors, as well as the values used for the current simulation."
 The CCD flatness techuologv aud testing is reviewed in Takacs et al. (, The CCD flatness technology and testing is reviewed in Takacs et al. (
2006) and Radcka et al. (,2006) and Radeka et al. (
2009).,2009).
 The focal plane errors that we assunie for the siuulatious in this paper are larger than the curreut expectation so that the results presented here are conservative., The focal plane errors that we assume for the simulations in this paper are larger than the current expectation so that the results presented here are conservative.
 Several veudors are now exceeding these flatuess specifications., Several vendors are now exceeding these flatness specifications.
 A realization of the LSST focal plane based ou these error distributions is displaved im Figure 7.., A realization of the LSST focal plane based on these error distributions is displayed in Figure \ref{fig_focal_plane}.
 The resulting behavior of the PSF ellipticity is illustrated in Figure 8.., The resulting behavior of the PSF ellipticity is illustrated in Figure \ref{fig_ellipticity_focal_plane_no_atmos}.
 These PSPs are obtained by evaluating optical path differeuce functions across the focal plane iu the abseuce of atmospheric turbulence using the ZEMAX software: see also Jarvis. Schechter. Jain (2008) for the discussion on the mipact of the telescope focus on the PSF ellipticity.," These PSFs are obtained by evaluating optical path difference functions across the focal plane in the absence of atmospheric turbulence using the ZEMAX software; see also Jarvis, Schechter, Jain (2008) for the discussion on the impact of the telescope focus on the PSF ellipticity."
 Several features are noteworthy in Figure 8.., Several features are noteworthy in Figure \ref{fig_ellipticity_focal_plane_no_atmos}.
 First. the aberrationdnuduced ellipticity is larec.," First, the aberration-induced ellipticity is large."
 Most of the large sticks in the plot exceed ~10% ellipticity. approaching rearly 30% at the field οσον.," Most of the large sticks in the plot exceed $\sim10$ ellipticity, approaching nearly $30$ at the field edges."
 Because the maxi deviation in this realization is siuall (~10 gran). these aree values of ellipticity remind us that the focal error olerauce of LSST is indeed narrow: however. we note hat in the ceutral region (S 17)) the scusitivity of PSF clougation to height error is somewhat mitigated.," Because the maximum deviation in this realization is small $\sim10$ $\mu$ m), these large values of ellipticity remind us that the focal error tolerance of LSST is indeed narrow; however, we note that in the central region $\lesssim 1$ ) the sensitivity of PSF elongation to height error is somewhat mitigated."
 Second. sharp discoutinuities of PSF ellipticity arc xeseut where the CCD heights also vary abruptly.," Second, sharp discontinuities of PSF ellipticity are present where the CCD heights also vary abruptly."
 Third. although it may bea bit difficult to recognize this feature in Fieure &.. a sinallscale variation is observed eve within a CCD iainly due to the tilt aud potato chi effects.," Third, although it may be a bit difficult to recognize this feature in Figure \ref{fig_ellipticity_focal_plane_no_atmos}, a small-scale variation is observed even within a CCD mainly due to the tilt and potato chip effects."
 The presence of these smooth. sinall-scale variation and discontinuous changes across CCD gaps are the most challenging aspects of the PSF description aud modeling for LSST.," The presence of these smooth, small-scale variation and discontinuous changes across CCD gaps are the most challenging aspects of the PSF description and modeling for LSST."
 Αν attempt to use a sinele set of polynomials to characterize the PSF behavior across the cutive focal plane fails because the simall-scale variation requires nupractical. high-order terms iuthe polynomials. aud in addition no interpolation scheme cam satisfactorily reproduce the sharp discoutimuitics across the CCD borders.," Any attempt to use a single set of polynomials to characterize the PSF behavior across the entire focal plane fails because the small-scale variation requires impractical, high-order terms inthe polynomials, and in addition no interpolation scheme can satisfactorily reproduce the sharp discontinuities across the CCD borders."
 This is the reason thatin the current paper we perform interpolation. CCD-by-CCD to model LSST PSFs., This is the reason thatin the current paper we perform interpolation CCD-by-CCD to model LSST PSFs.
 Considering the 1-2 cveles of the potato chip effect within a CCD. we estimate that 3-1 order polvuonials," Considering the 1-2 cycles of the potato chip effect within a CCD, we estimate that 3-4 order polynomials"
depending on the time the heat pulse is switched on.,depending on the time the heat pulse is switched on.
" As mentioned above, for a realistic situation, we have simulated the ignition of the whole loop system, with a gradually increasing number of simultancously heated strands."," As mentioned above, for a realistic situation, we have simulated the ignition of the whole loop system, with a gradually increasing number of simultaneously heated strands."
" The transient evolution will be the subject of a future work, but Fig."," The transient evolution will be the subject of a future work, but Fig."
 4 shows the initial evolution of the temperature averaged over the whole loop system., \ref{time_profile_T_med} shows the initial evolution of the temperature averaged over the whole loop system.
" In the following we will focus our attention to the final time (?=2000 s) when the loop system enters a presumably long steady state, in which the heat pulse ‘storm’ (1.c. the 2000 pulses in total) repeats continuously."," In the following we will focus our attention to the final time $t= 2000$ s) when the loop system enters a presumably long steady state, in which the heat pulse 'storm' (i.e. the 2000 pulses in total) repeats continuously."
" In à way, we are therefore implicitly assuming that the time taken to re-energise the magnetic field (c.g. by twisting, braiding) is about 2000 s. Fig."," In a way, we are therefore implicitly assuming that the time taken to re-energise the magnetic field (e.g. by twisting, braiding) is about 2000 s. Fig."
" 5 shows the distribution of the emission measure of the entire collection of strands, i.c. of the whole loop system, versus temperature (e.g. ?)) at this time."," \ref{EM_T_2000s} shows the distribution of the emission measure of the entire collection of strands, i.e. of the whole loop system, versus temperature (e.g. \citealt{Cargill_1994}) ) at this time."
" We show the distribution of the coronal part only, 1.c. upper of the loop bundle, excluding the lower layers."," We show the distribution of the coronal part only, i.e. upper of the loop bundle, excluding the lower layers."
" The peak of the distribution is around 3 MK, as planned."," The peak of the distribution is around 3 MK, as planned."
" We also see that the distribution 1s quite broad; the flatter tail is toward the cool part, but there are also significant hot components up to 10 MK, as expected, due to the presence of the strong heat pulses."," We also see that the distribution is quite broad; the flatter tail is toward the cool part, but there are also significant hot components up to 10 MK, as expected, due to the presence of the strong heat pulses."
" These hot components are of course minor, due to the small duty cycle of the heating phase with respect to the whole evolution of the strand."," These hot components are of course minor, due to the small duty cycle of the heating phase with respect to the whole evolution of the strand."
 Using the method described in the Sec 2.. we synthesize maps of emission for all the lines considered.," Using the method described in the Sec \ref{sec:model}, we synthesize maps of emission for all the lines considered."
 Fig., Fig.
" 6 (eft column) shows a subset of them, namely the map in cool EIS lines, i.c. Mg VIL O78 (logΤΙΚΙ= 5.8), medium temperature EIS line Fe X 2186 (logTILK]= 6.0). warm EIS line Fe XV 1284 (logT[K|= 6.4). and hot X-ray line Ca XIX 43.21 (logT[K|= 7.4)."," \ref{loop_profile} (left column) shows a subset of them, namely the map in cool EIS lines, i.e. Mg VII $\lambda 278$ $log ~ T[K] ~ = ~ 5.8$ ), medium temperature EIS line Fe X $\lambda 186$ $log ~ T[K] ~ = ~ 6.0$ ), warm EIS line Fe XV $\lambda 284$ $log ~ T[K] ~ = ~ 6.4$ ), and hot X-ray line Ca XIX $\lambda 3.21$ $log ~ T[K] ~ = ~ 7.4$ )."
" The images on the left column are analogous, and can be compared, to observed ones (e.g. ?)))."," The images on the left column are analogous, and can be compared, to observed ones (e.g. \cite{Tripathi_2009}) )."
 The grey scale is logarithmic and spans a factor 10 in all maps and plots., The grey scale is logarithmic and spans a factor $10$ in all maps and plots.
 This is a customary choice for showing observed data (c.g. ?)))., This is a customary choice for showing observed data (e.g. \cite{Tripathi_2009}) ).
 In the first three lines. the emission decreases from the base to the top of the loop. due to the density decreasing as well.," In the first three lines, the emission decreases from the base to the top of the loop, due to the density decreasing as well."
" We show the intensity of the upper 90% of the loop system, i.c. the coronal part only."," We show the intensity of the upper $90\%$ of the loop system, i.e. the coronal part only."
" In this way we exclude the emission from cromosphere and transition region, which, nevertheless. would appear saturated in this color"," In this way we exclude the emission from cromosphere and transition region, which, nevertheless, would appear saturated in this color"
Is the change of centroid. position in cconsistent with it undergoing mücrolensing at racio wavelengths. and if so. is this responsible for the observed absorption line variabilitv?,"Is the change of centroid position in consistent with it undergoing microlensing at radio wavelengths, and if so, is this responsible for the observed absorption line variability?"
 “Phe Solar mass Einstein radius in the source plane for is~S107h3per. and sources must be smaller than this if they are to be significantly. enhanced.," The Solar mass Einstein radius in the source plane for is $\sim8\times10^{-3}h^{-\frac{1}{2}}$, and sources must be smaller than this if they are to be significantly enhanced."
 Vhis corresponds to an angular scale of ~ 2microarcseconds., This corresponds to an angular scale of $\sim2$ microarcseconds.
 In their. study of the quadruple lens Q2237|0305. Lewis Ibata (1998) found that image shifts of 20-30 Einstein radii can occur on time scales substantially shorter than the crossing time of an Einstein racius.," In their study of the quadruple lens Q2237+0305, Lewis Ibata (1998) found that image shifts of 20-30 Einstein radii can occur on time scales substantially shorter than the crossing time of an Einstein radius."
 For this crossing time is 33/ vps. and the subsequent caustic crossing time will be of order weeks to months.," For this crossing time is $\sim33 h^{-\frac{1}{2}}$ yrs, and the subsequent caustic crossing time will be of order weeks to months."
 Hence. it is expected that the separation between the images in will change by ~LO100 microarescconds on this time scale.," Hence, it is expected that the separation between the images in will change by $\sim10-100$ microarcseconds on this time scale."
 Lt should be noted. however. that the degree of the expected astrometric shifts. are a function. of several parameters. especially the macrolensing shear.," It should be noted, however, that the degree of the expected astrometric shifts are a function of several parameters, especially the macrolensing shear."
 Given the ring-like nature of (Jauncy et al., Given the ring-like nature of (Jauncy et al.
 1991). this may be substantial in the vicinity of the images of the quasar core. potentially accounting for he large scale shifts observed by Garrett et al. (," 1991), this may be substantial in the vicinity of the images of the quasar core, potentially accounting for the large scale shifts observed by Garrett et al. ("
19907).,1997).
 More detailed. simulations. undertaken using the macrolensing xwameters for1830-211. are required before this can »e Lully addressed.," More detailed simulations, undertaken using the macrolensing parameters for, are required before this can be fully addressed."
 Η the observed line profile variability in is due to microlensing. then the clouds must be. of xuwsec/subparsec scales.," If the observed line profile variability in is due to microlensing, then the clouds must be of parsec/subparsec scales."
 Rather than rellecting individual clouds. however. these may represent. inhomogeneities in a arger scale absorption distribution. which could. potentially. »e fractal (IEbmegreen. 1997).," Rather than reflecting individual clouds, however, these may represent inhomogeneities in a larger scale absorption distribution, which could, potentially, be fractal (Elmegreen 1997)."
 Lt is apparent that is subject to gravitational microlensing to some degree. although. the evidence is currently only suggestive. that the observed. line profile. cillerences are due. to. the absorption mechanism outlined in. this paper.," It is apparent that is subject to gravitational microlensing to some degree, although the evidence is currently only suggestive that the observed line profile differences are due to the absorption mechanism outlined in this paper."
 Hence. further spectroscopic monitoring of the radio absorption lines. as. performed. by Wiklind. Combes (1998). is required.," Hence, further spectroscopic monitoring of the radio absorption lines, as performed by Wiklind Combes (1998), is required."
 Coupled with detailed. numerical simulations. such monitoring may provice clues to the distribution of absorbing material in galaxies on subparsec scales.," Coupled with detailed numerical simulations, such monitoring may provide clues to the distribution of absorbing material in galaxies on subparsec scales."
 Emploving a numerical approach. this paper has investigated the inlluence of à combination of microlensing ancl a distribution of absorbing material within the lensing galaxy upon temporal changes in the depth of absorption lines in quasar spectra.," Employing a numerical approach, this paper has investigated the influence of a combination of microlensing and a distribution of absorbing material within the lensing galaxy upon temporal changes in the depth of absorption lines in quasar spectra."
 Lt is found that the action of macrolensing compresses the large scale. absorption cloud information into smaller region. as seen by the source. whereas the microlensing ‘folds’ the absorption patten.," It is found that the action of macrolensing compresses the large scale absorption cloud information into smaller region, as seen by the source, whereas the microlensing `folds' the absorption patten."
 Several clouds clistributions were examined. and significant modulation of the depth of the absorption line resulted.," Several clouds distributions were examined, and significant modulation of the depth of the absorption line resulted."
 The form of the absorption line variability cdillerecl significantly from the simple case where the light from the quasar shone solely through an absorbing medium with no additional influence from gravitational lensing., The form of the absorption line variability differed significantly from the simple case where the light from the quasar shone solely through an absorbing medium with no additional influence from gravitational lensing.
 Hence. such variability could inlluence svstems sullering gravitational microlensing.," Hence, such variability could influence systems suffering gravitational microlensing."
 Due to the small size of the Einstein radius of stellar mass objects at cosmological distances. however. only sub-xwsec scale variations in absorbing material will result. in an observable effect.," Due to the small size of the Einstein radius of stellar mass objects at cosmological distances, however, only sub-parsec scale variations in absorbing material will result in an observable effect."
 The gravitationally lensecl quasar. possess several temporal features that are consistent. with it undergoing gravitational microlensing.," The gravitationally lensed quasar, possess several temporal features that are consistent with it undergoing gravitational microlensing."
 Furthermore. »possesses a number of prominent molecular absorption lines visible at radio wavelengths.," Furthermore, possesses a number of prominent molecular absorption lines visible at radio wavelengths."
 In one of these. HCO.(241. small. but significant variations in the absorption line profile over a period of six months have been reported.," In one of these, ${\rm HCO^+(2 \leftarrow 1)}$, small, but significant variations in the absorption line profile over a period of six months have been reported."
 Lhe analysis presented in this paper suggests that these changes may be due to the influence of gravitational microlensing., The analysis presented in this paper suggests that these changes may be due to the influence of gravitational microlensing.
 Lt is important to note. however. that currently very Little is known about the scale of structure of absorbing material," It is important to note, however, that currently very little is known about the scale of structure of absorbing material"
This simall value of Zt: obtained by minimizine tle width of GaiSsian fitted to SNe LF is an alternative manifestation of the low value of 3 obtained by Astieretal.(2006) by miniulsius the residual scatter in the Hubble diagram along with costYological pa‘ameters.,This small value of $R_V$ obtained by minimizing the width of Gaussian fitted to SNe LF is an alternative manifestation of the low value of $\beta$ obtained by \citet{SNLS} by minimising the residual scatter in the Hubble diagram along with cosmological parameters.
 Ixowalski aud Ixessleretal.(2009) inciated that minimizing he scatter iu he Hubble diagrar1 tends to eive Ry biased towards a value ower than the true valte., \citet{Kowalski08} and \citet{Kessler09} indicated that minimizing the scatter in the Hubble diagram tends to give $R_V$ biased towards a value lower than the true value.
 Orr suilatiou. however. shows that the value of 2 may be biasec to the lower value but no mo‘e thaiby ~J.Lo so that this caniol be the reason for 2y being sigiüficantly smaller han 3.1.," Our simulation, however, shows that the value of $\beta$ may be biased to the lower value but no more than by $\sim0.1$ , so that this cannot be the reason for $R_V$ being significantly smaller than 3.1."
 Wedo 1Ol coiclide here that Ry is actually smaller but take ο=1. Las our fiducial cloice. keeping iu nind he uncertainty [rom J?) iu the analysis in what follows.," Wedo not conclude here that $R_V$ is actually smaller but take $\beta=4.1$ as our fiducial choice, keeping in mind the uncertainty from $R_V$ in the analysis in what follows."
 I has been argued tiat the ruaximuimn uiminosity correlaes wlhi the ligit curve shape. or more specilically tie decline rate (Ας. stretch. o: SALT2s cry). aud the inclusion of the correlation with it makes the behaviour of the LF tielter (Pskovskii198[:Phillips1993:Hamuyetal.1996).," It has been argued that the maximum luminosity correlates with the light curve shape, or more specifically the decline rate $\Delta
m_{15}$, stretch, or SALT2's $x_1$ ), and the inclusion of the correlation with it makes the behaviour of the LF tighter \citep{Pskovskii84,Phillips93,Hamuy96}."
. We slow ln j»auels (c) aud (cl) of Figure 7 ile LF where briglituess is corrected by aay. with a chose 110 1uluunise the width of the Ciaussiau.," We show in panels (c) and (d) of Figure \ref{fig:snlf} the LF where brightness is corrected by $\alpha x_1$, with $\alpha$ chosen to minimise the width of the Gaussian."
 Iu pauel (c). 9 is fixed to our fiducial value of L1 auc tle minimisation glves a=0.052.," In panel (c), $\beta$ is fixed to our fiducial value of 4.1 and the minimisation gives $\alpha = 0.052$."
 In (d both a and 3 are chosen to minimise the width. which results in =2.52 amd a=0.123.," In (d) both $\alpha$ and $\beta$ are chosen to minimise the width, which results in $\beta = 2.52$ and $\alpha =
0.123$."
 The correlation of ry with brightness makes the Gaussia distribution larrower. especially for the case in which both a aud ;? are optiuised.," The correlation of $x_1$ with brightness makes the Gaussian distribution narrower, especially for the case in which both $\alpha$ and $\beta$ are optimised."
 The narrowest Gaussian is obtained with a=0.123 which corresponds to a’=0.76 where the correction for the light curve shape to b‘lightness is expressed as a’Amys., The narrowest Gaussian is obtained with $\alpha=0.123$ which corresponds to $\alpha^\prime=0.76$ where the correction for the light curve shape to brightness is expressed as $\alpha^\prime \Delta m_{15}$.
 This value is consjistent. with 7820.18 obtained by Hamuyetal.(1996)., This value is consistent with $\pm0.18$ obtained by \cite{Hamuy96}.
. Figure 8. shows tle correlation between brightness of SNe with the shaye parameter aid the coloir excess., Figure \ref{fig:Mab} shows the correlation between brightness of SNe with the shape parameter and the colour excess.
 This correlation is the reason. why the width of the luminosity function decreases upon inclusious of the ight curve shape parameter aud the colour excess pa‘ammeter (with a sinall Ih)., This correlation is the reason why the width of the luminosity function decreases upon inclusions of the light curve shape parameter and the colour excess parameter (with a small $R_V$ ).
" The correlation ii the upper panel is represented by Aj;=—19.312-2.93c as shown in Figure T((b). aud that in the lower pauel by Alp—ble=—19.130.052:r, as in Fieure το)."," The correlation in the upper panel is represented by $M_B=-19.34+2.93c$ as shown in Figure \ref{fig:snlf}( (b), and that in the lower panel by $M_B-4.1c=-19.43-0.052x_1$ as in Figure \ref{fig:snlf}( (c)."
 The LEs of SNe Ia host galaxies. as calciated by the 1/Vinay method. are shown in Figu‘es O.. 10 aud LL with solid histograms for the r. g. a1da z-passbauds.," The LFs of SNe Ia host galaxies, as calculated by the $1/V_{\rm max}$ method, are shown in Figures \ref{fig:hostlf_r}, \ref{fig:hostlf_g}
 and \ref{fig:hostlf_i} with solid histograms for the $r$, $g$, and $i$ -passbands."
 Hostless SNe are indicated by 5iacled istograun at rightimnost bin., Hostless SNe are indicated by shaded histogram at rightmost bin.
 We draw with solk| curves the LF of general field galaxies obtained by lautonetal.(2001).. uultiplied with the Itinosity.," We draw with solid curves the LF of general field galaxies obtained by \citet{Blanton01}, multiplied with the luminosity."
 The curves slow good match of the LF of : host galaxies with that of the field galaxies. Which means that the LF of galaxies derive from e Ta faithfully represents that of galaxies | the field.," The curves show good match of the LF of SN host galaxies with that of the field galaxies, which means that the LF of galaxies derived from SNe Ia faithfully represents that of galaxies in the field."
 We do uot see any particular deviations between the two for three colour passbauds g. rand we studied. meaning that the occuJTVence [9] “SNe Ia is primarily proportional to the lu1inosity of galaxies.," We do not see any particular deviations between the two for three colour passbands$g$ , $r$ and $i$ we studied, meaning that the occurrence of SNe Ia is primarily proportional to the luminosity of galaxies."
" Matching the two LE's gives the 5e Ia rate in the conventional superuova uit (SNu). the SN rate per 10!"" solar luminosiy per"," Matching the two LF's gives the SNe Ia rate in the conventional supernova unit (SNu), the SN rate per $10^{10}$ solar luminosity per"
structures towards high redshift radio sources showing high seusitivitv to eas temperature (Nuetal.2009).,structures towards high redshift radio sources showing high sensitivity to gas temperature \citep{XuC09}.
. The problem of the 21 cm forest signatures produced bw different Xiuds of structures has been addressed by several autliors., The problem of the 21 cm forest signatures produced by different kinds of structures has been addressed by several authors.
 Carillietal.(2002) presented a detailed study of 21 cin absorption bv the mean joutra IGM as well as filamentary structures based on the shuulatious of Cinediun.(2000)... but their box was too small to account for large scale structures and was not able to resolve collapsed: objects.," \citet{Carilli02} presented a detailed study of 21 cm absorption by the mean neutral IGM as well as filamentary structures based on the simulations of \citet{Gnedin00}, but their box was too small to account for large scale structures and was not able to resolve collapsed objects."
 Tustead. Furlauctto&Loeb(2002) used a simple analytic model to compute the absorption profiles and abundances of minibhalos aud galactic disks.," Instead, \citet{Furlanetto02} used a simple analytic model to compute the absorption profiles and abundances of minihalos and galactic disks."
 Later on. Fulanetto(2006) re-exanuded four kiuds of 21 cin forest absorption signatures iu a broader context. especially the transniüssion gaps produced by ionized. bubbles.," Later on, \citet{Furlanetto06} re-examined four kinds of 21 cm forest absorption signatures in a broader context, especially the transmission gaps produced by ionized bubbles."
 Receuthy Xuetal.(2010) eveloped a more detailed model of the 21 cmi absorption lines of unimihalos (ie. starless galaxies. MITIS) and dwart galaxies (star-forming galaxies. DCs) during the epoch of reionization. explored the physical origins of the line profiles. aud geucrated svuthetic spectra of the 21 cia forest on top of both hieli-- quasars and ezaumna rav burst (GRB) afterelows.," Recently, \citet{XuF10} developed a more detailed model of the 21 cm absorption lines of minihalos (i.e. starless galaxies, MHs) and dwarf galaxies (star-forming galaxies, DGs) during the epoch of reionization, explored the physical origins of the line profiles, and generated synthetic spectra of the 21 cm forest on top of both $z$ quasars and gamma ray burst (GRB) afterglows."
 Tuterestinely. they fud that: (1) MIIS and DCs show very distinct 21 cimi line absorption profiles ai) rev contribute differentle to the spectra. due to the nass segregation between the two populations.," Interestingly, they find that: (i) MHs and DGs show very distinct 21 cm line absorption profiles (ii) they contribute differently to the spectrum due to the mass segregation between the two populations."
 It follows that it is in principle possible to build a criterion based ou the 21 cm forest προςπα to efficiently select. DCs., It follows that it is in principle possible to build a criterion based on the 21 cm forest spectrum to efficiently select DGs.
 The goal of this work is to derive the different sjeuatures of DCs and MIIS using a 21 ci spectrmm of liel-: radio sources. and provide a criterion to pick DCs lines in the spectra.," The goal of this work is to derive the different signatures of DGs and MHs using a 21 cm spectrum of $z$ radio sources, and provide a criterion to pick DGs lines in the spectrum."
 For jose candiates. precise redshift iformation will ο ανααυoe: moreover. given tfje angular TOSion of the backeround source. he 21 cm ‘Orest observatikon provides an excelleut tool for Ocaπιο the rz. DCs.," For these candidates, precise redshift information will be available; moreover, given the angular position of the background source, the 21 cm forest observation provides an excellent tool for locating the $z$ DGs."
 The exeat acvantage of using high-: CRBs as background radio sources is hat the followap IR JWST observatioi after the afterelow has faded away will not be hampered bx he presence of a very luniuous source (as in the ease of a backeround quasar) in thefield!., The great advantage of using $z$ GRBs as background radio sources is that the follow-up IR JWST observations after the afterglow has faded away will not be hampered by the presence of a very luminous source (as in the case of a background quasar) in the.
. Tere we briefly summarize the main features of the model used in this work. but refer the interested reader to Nuetal.(2010). for a more comprehensive description.," Here we briefly summarize the main features of the model used in this work, but refer the interested reader to \citet{XuF10} for a more comprehensive description."
" We mse the Tormen halo mass function (Sheth&Tormen1999) to model the halo number deusity at high redshift iu the mass range LO?29A, which covers the mininuun mass allowed to collapse (Abeletal.2000:O'Shea&Norinan2007) and most of the galaxies that are responsible for reiouization (Choudhury&Ferrara2007)."," We use the Sheth-Tormen halo mass function \citep{ST99} to model the halo number density at high redshift in the mass range $10^{5-10}M_\odot$, which covers the minimum mass allowed to collapse \citep{Abel00,ON07} and most of the galaxies that are responsible for reionization \citep{CF07}."
". The dark matter halos have an NEW density profile within the viral radii sq (Navarro.Frenk&White1997).. with a concentration parameter fitted to ligh- simulation results by Ciaoetal.(200511: the dark iatter density and velocity. structure outside⋅ r4, are described⋅ by an ""Iufall. (Barkana2001)."," The dark matter halos have an NFW density profile within the virial radii $r_{\rm vir}$ \citep{NFW97}, with a concentration parameter fitted to high-redshift simulation results by \cite{Gao05}; the dark matter density and velocity structure outside $r_{\rm vir}$ are described by an “Infall \citep{Barkana04}."
".. The easons duside the ma, Is assuned to be in hwdrostatie equilibrium at temperature Zi, in the dark matter potential. while the seas outside follows the dark matter overdeusitv aud velocity field."," The gas inside the $r_{\rm vir}$ is assumed to be in hydrostatic equilibrium at temperature $T_{\rm vir}$ in the dark matter potential, while the gas outside follows the dark matter overdensity and velocity field."
 Ouce the halo population is fixed. a timescale criterion for star formation ds introduced to determine whether a halo is capable of hosting star formation.," Once the halo population is fixed, a timescale criterion for star formation is introduced to determine whether a halo is capable of hosting star formation."
" The timescale required for turing on star formation is modeled as the masiuuun )etween the free-fal time τμ aud the II» cooling nie feo, (Teemarkoetal.1997).. Ίνα, fup—nanfe. tesa}."," The timescale required for turning on star formation is modeled as the maximum between the free-fall time $t_{\rm ff}$ and the $_2$ cooling time $t_{\rm cool}$ \citep{Tegmark97}, i.e. $t_{\rm SB} = \max \{\, t_{\rm ff},\, t_{\rm
cool}\,\}$ ."
 Then star formation activitv ⋝↸∖∶↴∙⊾↕∐↴∖↴⋜↧↑⊺∠∖∶⊺∠⇪↖⊺∠⊨↴⊔∙↖↖↽∐↸∖↥⋅↸∖⊺∠⇪↕↴∖↴↑↕∐∖∐⋜↧↕∪ ≯∪↥⋅∐↓⋜↧↑↕∪∐↑↕⊔↸∖↻↥⋅↸∖≺∐≼⊳↑↸∖≺↧↴⋝∙↖⇁↑∐↸∖↴∖↴↑⋜⋯≼↧⋜∐⋅≼↧⊏↕⋟≋ ⊔∪≼∐∖↕∏⇀⋜↧↸⊳↸∖⋅↖⇁∙∖↽≼⊲∪↕↸∖↕∩," Then star formation activity begins at $t_{\rm s} =
t_{\rm F} + t_{\rm SB}$, where $t_{\rm F}$ is the halo formation time predicted by the standard EPS model \citep{LC93}."
⋂∶≩⋝⋈↕↕≯∤∖↕↴∖↴↕⋜∐⋅∶↴∙⊾↸∖↥⋅↑∐⋜⋯ he IInbble nue at the halo redshift. we define he system as aginihalo. o. a starless salaxx.," If $t_{\rm
s}$ is larger than the Hubble time at the halo redshift, we define the system as a, i.e. a starless galaxy."
 The ionized fraction in a MIT is computed with collisional ionization equilibria. which depends onu is temperatire.," The ionized fraction in a MH is computed with collisional ionization equilibrium, which depends on its temperature."
" The gas within rg Is at the virial temperature. and in the absence of au κ. background the σας outside is adiabatically conrpressed. so that the temperature is simply Tu=Τον|8Ó) Ἡ, where 5=5/3 is the"," The gas within $r_{\rm vir}$ is at the virial temperature, and in the absence of an X-ray background the gas outside is adiabatically compressed, so that the temperature is simply $T_{\rm K} = T_{\rm IGM}
(1+\delta)^{\gamma-1}$ , where $\gamma = 5/3$ is the"
photous is at the order of ~10 oe 7. which is comparable with the maeuectic energy density for a field strength of ~SC. Therefore the IC power cannot be ignored.,"photons is at the order of $\sim10^{-12}$ erg $^{-3}$, which is comparable with the magnetic energy density for a field strength of $\sim5\mu$ G. Therefore the IC power cannot be ignored."
 Iu the IC scenario. he required. Loreutz factor to produce ~10 keV yhotous is only at the order of ~104 which is uuchn easier to achieve.," In the IC scenario, the required Lorentz factor to produce $\sim10$ keV photons is only at the order of $\gamma\sim10^{4}$ which is much easier to achieve."
 The major. difficulty⋅⋅ for: explainingDa the N-rav- ict of jiu the context of PWN comes from its orientation with respect to the pulsus proper motion., The major difficulty for explaining the X-ray jet of in the context of PWN comes from its orientation with respect to the pulsar's proper motion.
 The jet protrudes from the pulsar aud extends toward jorthweest. which makes an augle of ~118? from he proper motion direction of the pulsar (cf.," The jet protrudes from the pulsar and extends toward northwest, which makes an angle of $\sim118^{\circ}$ from the proper motion direction of the pulsar (cf."
 Fie. 1))., Fig. \ref{xmm_img}) ).
 This peculiar orientation makes the bow-shock interpretation questionable as this scenario requires the nebular emission lies iu au opposite direction of the pulsar’s velocity vector., This peculiar orientation makes the bow-shock interpretation questionable as this scenario requires the nebular emission lies in an opposite direction of the pulsar's velocity vector.
 Iu view of this difficulty. different explanations for the nature of the jet association with have been proposed. ranging from the maeguetically confinement of relativistic particles leakius from the bow-shock (Bauclicra 2008) to a picture similar to the AGN jet outflow (Johnson Wane 2010).," In view of this difficulty, different explanations for the nature of the jet association with have been proposed, ranging from the magnetically confinement of relativistic particles leaking from the bow-shock (Bandiera 2008) to a picture similar to the AGN jet outflow (Johnson Wang 2010)."
 Nevertheless. there is no couseusus ou its plivsical origin vet.," Nevertheless, there is no consensus on its physical origin yet."
 Since the extended feature of PSB J0357|3205 is akin to that of65.. a comparative analvsis between these two systems nieht provide us a deeper insight on the nature of their jets.," Since the extended feature of PSR J0357+3205 is akin to that of, a comparative analysis between these two systems might provide us a deeper insight on the nature of their jets."
 While the distance aud the proper motion direction of are well-coustrained. these are the most important paralucters to be determinect in further investieatiou of PSR J0357|3205.," While the distance and the proper motion direction of are well-constrained, these are the most important parameters to be determined in further investigation of PSR J0357+3205."
 These would enable a more reliable estimations of its cuerectic and most iuportautlv the orientation of the jet with respect to the pulsas space velocity. which will allow us to ascertain if these two svstenis have a simular physical origin.," These would enable a more reliable estimations of its energetic and most importantly the orientation of the jet with respect to the pulsar's space velocity, which will allow us to ascertain if these two systems have a similar physical origin."
 Hs certainly one of the most spectacular neutron stars regarding its motion through space as its transverse velocity. ο=865(d/1kpc} lau ς lis among the highest in the curently known pulsar population.," is certainly one of the most spectacular neutron stars regarding its motion through space as its transverse velocity, $v=865(d/1~{\rm kpc})$ km $^{-1}$, is among the highest in the currently known pulsar population."
 Ou the other laud. the well measured bow shock with its inclination iu the plane of the sky sugeests a neelieible radial velocity componucut (Chatterjee Cordes 2001) which has Όσοι confirmed recently by Tetzlaff ct al. (," On the other hand, the well measured bow shock with its inclination in the plane of the sky suggests a negligible radial velocity component (Chatterjee Cordes 2004) which has been confirmed recently by Tetzlaff et al. ("
2009) who sugeest the origin of tto lie within the Cyeuus OD3 association or a ΠΕ massive cluster nearby.,2009) who suggest the origin of to lie within the Cygnus OB3 association or a small massive cluster nearby.
" This origi indicates a ounsmall radialsoial ovvelocityJp ofot e,—5gl5l21/2, kins 1.", This origin indicates a small radial velocity of $v_{r}=-21^{+81}_{-70}$ km $^{-1}$.
 The proposed kincmatic age of των20.8 Myr is consistent with a characteristic pulsar age of Τομ=LL Myr (Hobbs et al., The proposed kinematic age of $\tau_{\rm kin}\approx 0.8$ Myr is consistent with a characteristic pulsar age of $\tau_{\rm char}=1.1$ Myr (Hobbs et al.
 200.b Manchester et al., 2004; Manchester et al.
 2005)., 2005).
 To advance the analvsis we preseuted earlier (Tetzlaff et al., To advance the analysis we presented earlier (Tetzlaff et al.
 2009). we attempt to identify a runaway star which aight have been the former companion of this pulsus progenitor star (we refer to such runway stars as binary supernova scenario (BSS) ruuawars: Blaamy 1961).," 2009), we attempt to identify a runaway star which might have been the former companion of this pulsar's progenitor star (we refer to such runaway stars as binary supernova scenario (BSS) runaways; Blaauw 1961)."
 Therefore. we take all runway star candidates with full kinematics from the rumaway star catalogue of Tetzlaff et al. (," Therefore, we take all runaway star candidates with full kinematics from the runaway star catalogue of Tetzlaff et al. ("
2011) and compare the past flight oitlis of aand cach ruuaway star varving the observables by verformineg a Monte-Carlo simulation to fiud close encounters between the pulsu aud the runaway star.,2011) and compare the past flight paths of and each runaway star varying the observables by performing a Monte-Carlo simulation to find close encounters between the pulsar and the runaway star.
 Suuaultaneouslve we caleulate the distauce Oo possible parent associations/clusters listed in Tetzlaff ct al. (," Simultaneously, we calculate the distance to possible parent associations/clusters listed in Tetzlaff et al. ("
2010) (for a full description of the xocedure we refer to Tetzlaff et al.,2010) (for a full description of the procedure we refer to Tetzlaff et al.
 2009 and Tetzlaff et al., 2009 and Tetzlaff et al.
" For wwe take the following data for the equatorial coordinates. right ascension o aud declination à. the distance to the Sun d aud the proper motion conmonents pf) (5,=pn 0080) aud prs (Harrison et al."," For we take the following data for the equatorial coordinates, right ascension $\alpha$ and declination $\delta$, the distance to the Sun $d$ and the proper motion components $\mu_{\alpha}^*$ $\mu_{\alpha}^{*}=\mu_{\alpha}\cos\delta$ ) and $\mu_{\delta}$ (Harrison et al."
 1993: Taylor Cordes 1993: Yuan et al., 1993; Taylor Cordes 1993; Yuan et al.
 2010):, 2010);
applying a statistical correction for dust extinction (222) orby combining observations in UV and with those in the IR wavelength range (?)..,"applying a statistical correction for dust extinction \citep{Kennicutt1983,Calzetti1994,Calzetti2000} or by combining observations in UV and with those in the IR wavelength range \citep{Kennicutt2009}."
 In order (ο estimate the SFR from CC SN rate measurements we have to assume the mass range of CC SN progenitors and to correct the rates for the fraction of the extinguished CC SNe that are missed in optical searches., In order to estimate the SFR from CC SN rate measurements we have to assume the mass range of CC SN progenitors and to correct the rates for the fraction of the extinguished CC SNe that are missed in optical searches.
 The lower mass limit for CC SN progenitors from direct detections of progenitor stars in high-resolution images has arrived at a best estimate of 8.57}5 (9)... which is in reasonable agreement with the most massive white dwarf progenitors(??)..," The lower mass limit for CC SN progenitors from direct detections of progenitor stars in high-resolution images has arrived at a best estimate of $8.5^{+1}_{-1.5}$ \citep{Smartt2009}, which is in reasonable agreement with the most massive white dwarf progenitors\citep{Williams2007,Williams2009}."
 This has led ?. to suggest that the current best estimate from these two methods is 8+|msun.," This has led \citet{2009ARA&A..47...63S}
 to suggest that the current best estimate from these two methods is $8\pm1$."
. If we assume this value of the observed CC SN rates in the galaxy samples A. B and C imply SFRs are plotted in Fig 6..," If we assume this value of the observed CC SN rates in the galaxy samples A, B and C imply SFRs are plotted in Fig \ref{sfrcom}."
 The observed CC SN rate ts of course only a robust lower limit since we have not applied any correction for undetected SNe., The observed CC SN rate is of course only a robust lower limit since we have not applied any correction for undetected SNe.
 The SFR from CC SNe is higher by a factor two compared with those in Sample A and C based on while there is good agreement with SFR based on that suggests we are not missing a large number of CC SNe within MMpe due to dust extinction. intrinsically faint magnitudes. or over-estimating the control time.," The SFR from CC SNe is higher by a factor two compared with those in Sample A and C based on while there is good agreement with SFR based on that suggests we are not missing a large number of CC SNe within Mpc due to dust extinction, intrinsically faint magnitudes, or over-estimating the control time."
 The main source of the difference between the corrected SFRs based on and is due likely to the attenuation corrections., The main source of the difference between the corrected SFRs based on and is due likely to the attenuation corrections.
 A few of the galaxies with the highest SFRs tend to show especially large discrepancies between and derived SFRs. and we suspect that some of these may arise from spurious causes such as extremely heavy extinction in edge-on systems (e.g.. M82). very large foreground Galactic extinction (e.g.. NGC 6946). or poorly measured fluxes (e.g... NGC 6744).," A few of the galaxies with the highest SFRs tend to show especially large discrepancies between and derived SFRs, and we suspect that some of these may arise from spurious causes such as extremely heavy extinction in edge-on systems (e.g., M82), very large foreground Galactic extinction (e.g., NGC 6946), or poorly measured fluxes (e.g., NGC 6744)."
 These galaxies carry disproportionate weight in the total SFRs for the samples. but even taking them into account the based SERs remain systematically larger.," These galaxies carry disproportionate weight in the total SFRs for the samples, but even taking them into account the based SFRs remain systematically larger."
 A small part of the offset comes from the adoption of the ? formula for estimating extinction corrections., A small part of the offset comes from the adoption of the \cite{Buat2005} formula for estimating extinction corrections.
 ? compared attenuations derived from that method with those from .t+TIR and +240 μι schemes. and found that the former are systematically larger. by about mmag.," \citet{Kennicutt2009} compared attenuations derived from that method with those from +TIR and +24 $\mu$ m schemes, and found that the former are systematically larger, by about mag."
 There is also a more important systematic offset (30-40%)) in TIR luminosity between MIPS. which was used for nearly all of our sample. and IRAS. which was used to calibrate the ? relation (formoredetailsseeFigures1-2of ?)..," There is also a more important systematic offset ) in TIR luminosity between MIPS, which was used for nearly all of our sample, and IRAS, which was used to calibrate the \cite{Buat2005} relation \cite[for more details see Figures 1-2 of][]{Kennicutt2009}."
 This difference 1s only important for galaxies with cold IRAS colours (where basically the IRAS wavelength coverage is not sufficient to integrate the IR emission reliably)., This difference is only important for galaxies with cold IRAS colours (where basically the IRAS wavelength coverage is not sufficient to integrate the IR emission reliably).
 Unfortuately that colour regime applies to most of our galaxy sample., Unfortunately that colour regime applies to most of our galaxy sample.
 The comparison between CC SN rate and SER based on other diagnosties has also bee done at larger volumes (2???) and points out à discrepancy in the opposite direction. with respect to the local Universe since the observed CC SN rate is lower than the predicted oe from SER measurements.," The comparison between CC SN rate and SFR based on other diagnostics has also been done at larger volumes \citep{Dahlen2004,Botticella2008,Bazin2009,Horiuchi2011} and points out a discrepancy in the opposite direction with respect to the local Universe since the observed CC SN rate is lower than the predicted one from SFR measurements."
 It is interesting to note that this discrepancy (about a factor two) is constant in. a large range of redshift (??)..," It is interesting to note that this discrepancy (about a factor two) is constant in a large range of redshift \citep{Botticella2008,Horiuchi2011}. ."
 9 estimated the SFR density in four different redshift bins, \cite{Dale2010} estimated the SFR density in four different redshift bins
0.3 auc 0.5 are probably members of the cluster.,0.3 and 0.5 are probably members of the cluster.
 Roughly 1/3 (3856) of our emission line stars belong to this group., Roughly 1/3 $38\%$ ) of our emission line stars belong to this group.
 Stars with E(B—V)>0.5 could be background objects., Stars with $E(B-V)>0.5$ could be background objects.
 Roughly LOY of the emission line stars has larger reddening than the cluster members., Roughly $40\%$ of the emission line stars has larger reddening than the cluster members.
" According to the .A,-distauce relation of Neckeletal.(1980).. these stars have cistauces larger than 1.9 kpe."," According to the $A_v$ -distance relation of \citet{neck80}, these stars have distances larger than 1.9 kpc."
 The distance cau be as large as 6.6 kpc lor the highly reddened O star id2905 (HD226868)., The distance can be as large as 6.6 kpc for the highly reddened O star id2905 (HD226868).
 Such large distances seem iniprobable. because they imply unreasonably large absolute magnitudes.," Such large distances seem improbable, because they imply unreasonably large absolute magnitudes."
 The calculated z value above the galactic plaue (z>200 pc) Lor id2905 also coutrasts with the accepted 2 values for ο - B stars (50—100 pc). which suggests that this star has extra circumstellar recdeniug compared to cluster stars.," The calculated $z$ value above the galactic plane $ z > 200$ pc) for id2905 also contrasts with the accepted $z$ values for O - B stars $50 - 100$ pc), which suggests that this star has extra circumstellar reddening compared to cluster stars."
 The distances of the other background stars seem reasonable for the position of NGC 6871 but some of these stars might be cluster members with extra reddening due to ciretumstellar material., The distances of the other background stars seem reasonable for the position of NGC 6871 but some of these stars might be cluster members with extra reddening due to circumstellar material.
 We examine this possibility further in the uext subsection., We examine this possibility further in the next subsection.
" Having derived the B-V color excess for the emission line stars. we calculate the 4, ancl the bolometric correction lor each star in our sample aud coustruct the HRD for each subsample specified iu the previous section (foreerouucl. cluster. backerouud)."," Having derived the B-V color excess for the emission line stars, we calculate the $A_v$ and the bolometric correction for each star in our sample and construct the HRD for each subsample specified in the previous section (foreground, cluster, background)."
 Fig 3 shows these ciagranms., Fig 3 shows these diagrams.
 Fig 3a (top left) shows the HRD for all stars in our sample. Fig 3b (top right) shows the HRD for the cluster: Fig 3c and Fig 3d (bottom left aud right) show the HRD of the background aud the foreground stars respectively.," Fig 3a (top left) shows the HRD for all stars in our sample, Fig 3b (top right) shows the HRD for the cluster; Fig 3c and Fig 3d (bottom left and right) show the HRD of the background and the foreground stars respectively."
 Small grey. dots denote normal main sequence stars: large black dots denote emission line stars., Small grey dots denote normal main sequence stars; large black dots denote emission line stars.
 The solid curves show the ZAMS of Siessetal.(2000). shifted with the distance modulus DAL=11.08 which corresponds to a distauce of 1619 pe 10901)., The solid curves show the ZAMS of \citet{siess00} shifted with the distance modulus $DM=11.08$ which corresponds to a distance of 1649 pc \citep{bat91}.
 Although the majority of emission line stars in the cluster follow the ZAMS. four emission line stars lie above the main sequence.," Although the majority of emission line stars in the cluster follow the ZAMS, four emission line stars lie above the main sequence."
 These objects may be PMS stars (see section [.)., These objects may be PMS stars (see section 4.).
 Several other nou-emission line objects lie close to these stars in the HRD., Several other non-emission line objects lie close to these stars in the HRD.
 The relatively large scatter of these apparent PMS stars in the HRD suggest a large age-spread in tlie cluster 1989).," The relatively large scatter of these apparent PMS stars in the HRD suggest a large age-spread in the cluster \citep{mass95,reim89}."
. The HRD of background objects (Fig 3c) is very similar to the HRD for cluster members., The HRD of background objects (Fig 3c) is very similar to the HRD for cluster members.
 Roughly half of the emission line stars lie on or above the ZAMS. which suggestsMOD that they are cluster members with extra reddeuing due to circumstellar material.," Roughly half of the emission line stars lie on or above the ZAMS, which suggests that they are cluster members with extra reddening due to circumstellar material."
 Radial velocity measuremenW. would test this possibilitv., Radial velocity measurements would test this possibility.
 The stars with lower reddening occupy au area above the main sequeuce., The stars with lower reddening occupy an area above the main sequence.
 Most of these stars are [oreerouud objects with different distance moduli., Most of these stars are foreground objects with different distance moduli.
 Three of the four zero reddening stars may lie at the same distance: they are on a line nearly parallel to the ZAMS aud may be a stall group oL young stars., Three of the four zero reddening stars may lie at the same distance: they are on a line nearly parallel to the ZAMS and may be a small group of young stars.
 The fourth of the unreddened stars may lie within 100 pe but its spectral type is uncertain., The fourth of the unreddened stars may lie within 100 pc but its spectral type is uncertain.
 Further observations would place better limits ou its distance., Further observations would place better limits on its distance.
 The other foreground stars lie just above the main sequence., The other foreground stars lie just above the main sequence.
 Their position is consistent with the clistance calculated, Their position is consistent with the distance calculated
Figure 1 shows the average velocity of the H53a line.,Figure 1 shows the average velocity of the $\alpha$ line.
 The velocity gradient is approximately 8 kms! between 57 and 65 kms! VLSR with a position angle —30° (west of North)., The velocity gradient is approximately 8 $^{-1}$ between 57 and 65 $^{-1}$ VLSR with a position angle $-30^\circ$ (west of North).
 This velocity pattern is consistent with either rotation or outflow., This velocity pattern is consistent with either rotation or outflow.
 We prefer the interpretation that the H53a velocities indicate rotation for a couple of reasons., We prefer the interpretation that the $\alpha$ velocities indicate rotation for a couple of reasons.
" First, the observed velocity gradient is perpendicular to the CO outflow."," First, the observed velocity gradient is perpendicular to the CO outflow."
 Bipolar outflows associated with lower mass stars are always oriented approximately perpendicular to the rotation., Bipolar outflows associated with lower mass stars are always oriented approximately perpendicular to the rotation.
" Second, the interpretation of rotation is consistent with the magnitude and direction of the rotational velocity gradient observed in NH3 and CH3CN (Zhang&Ho1997;Zhangetal.1998)."," Second, the interpretation of rotation is consistent with the magnitude and direction of the rotational velocity gradient observed in $_3$ and $_3$ CN \citep{ZhangHo1997, ZhangHoOhashi1998}."
". The clearest signature of rotation in the molecular gas is in figure 7 of in the position-velocity diagram of NH3(3,3) along a line at position angle of 135°, and consistent (135 - 180 = —55°) with the position-angle of the velocity gradient in the ionized gas (—30°)."," The clearest signature of rotation in the molecular gas is in figure 7 of \citet{ZhangHo1997} in the position-velocity diagram of $_3$ (3,3) along a line at position angle of $^\circ$, and consistent (135 - 180 =$-55^\circ$ ) with the position-angle of the velocity gradient in the ionized gas $-30^\circ$ )."
" We prefer to use the NH3(3,3) line in this interpretation rather than the (2,2) line because the higher excitation line should derive from gas closer to the hot HII region."," We prefer to use the $_3$ (3,3) line in this interpretation rather than the (2,2) line because the higher excitation line should derive from gas closer to the hot HII region."
" However, there is considerable uncertainty in locating the exact angle of the axis of rotation using the molecular line observations."," However, there is considerable uncertainty in locating the exact angle of the axis of rotation using the molecular line observations."
 These observations map the flow at relatively large scales before the accretion has spun-up by angular momentum conservation to higher rotational velocities., These observations map the flow at relatively large scales before the accretion has spun-up by angular momentum conservation to higher rotational velocities.
" The observed infall velocities are comparable or greater than the rotational velocities, and the rotational component is difficult to extract."," The observed infall velocities are comparable or greater than the rotational velocities, and the rotational component is difficult to extract."
" For example, comparison of the position angle (-55?) of the rotational gradient derived from the NH; observations with the position angle (—70°) derived from CH3CN (Zhangetal.1998) suggests an uncertainty of several tens of degrees."," For example, comparison of the position angle $^\circ$ ) of the rotational gradient derived from the $_3$ observations with the position angle $-70^\circ$ ) derived from $_3$ CN \citep{ZhangHoOhashi1998} suggests an uncertainty of several tens of degrees."
" Our observations do not fully resolve the HII region, but the 1.3 cm continuum observation of Keto,Zhang&Kurtz(2008) indicates a FWHM of 0.4"" similar to the upper limit estimated by Zhang&Ho(1997) and Scott(1978)."," Our observations do not fully resolve the HII region, but the 1.3 cm continuum observation of \citet{KetoZhangKurtz2008} indicates a FWHM of $0.4^{\prime\prime}$ similar to the upper limit estimated by \citet{ZhangHo1997} and \citet{Scott1978}."
". The radius, R, associated with the velocity difference AV, is uncertain because of our low angular resolution."," The radius, $R$, associated with the velocity difference $\Delta V$, is uncertain because of our low angular resolution."
 So we use the FWHM of the continuum emission as a characteristic size., So we use the FWHM of the continuum emission as a characteristic size.
" Assuming that the observed velocities are rotational, the mass of the star within the W51e2 region is estimated as M=AV?R/2G."," Assuming that the observed velocities are rotational, the mass of the star within the W51e2 region is estimated as $M = \Delta V^2R/2G $."
" Assuming a distance of 8 kpc, we derive a lower limit to the velocity gradient of >500 kms! pc! and a dynamical mass of V?R/2G>15 Ms This mass is roughly consistent with a previous estimate of about 20 M, derived from the emission measure of the radio continuum (Scott 1978).."," Assuming a distance of 8 kpc, we derive a lower limit to the velocity gradient of $> 500$ $^{-1}$ $^{-1}$ and a dynamical mass of $V^2R/2G > 15$ $_\odot$ This mass is roughly consistent with a previous estimate of about 20 $_\odot$ derived from the emission measure of the radio continuum \citep{Scott1978}. ."
 This estimated velocity gradient in the ionized gas is similar to that measured in molecular gas., This estimated velocity gradient in the ionized gas is similar to that measured in molecular gas.
 Zhang&Ho(1997) measure a velocity gradient in NH3 of 500 kms! pc! , \citet{ZhangHo1997} measure a velocity gradient in $_3$ of 500 $^{-1}$ $^{-1}$ 
rale Qe<3 107 + (corresponding to 3 ke + in water) was placed spectroscopically (Hsieh et al.,rate $Q_{CN} \le$ $\times$ $^{23}$ $^{-1}$ (corresponding to 3 kg $^{-1}$ in water) was placed spectroscopically (Hsieh et al.
 90110)., 2011c).
 Hsieh et al. (, Hsieh et al. (
20110) show (neglecting possible non-eravitational forces due to outgassing) that the orbit of P/2010 R2 is stable on timescales «1060 Myr and argue that this object was likely formed in-situ.,2011c) show (neglecting possible non-gravitational forces due to outgassing) that the orbit of P/2010 R2 is stable on timescales $\sim$ 100 Myr and argue that this object was likely formed in-situ.
 Fernandez οἱ al. (, Fernandez et al. (
1997) analvzed (wo photographic plates taken on 1949 November 19. when at A8 = L143 AU.,"1997) analyzed two photographic plates taken on 1949 November 19, when at $R$ = 1.148 AU."
 The object is trailed in both but. in the blue plate. shows a prominent dilfuse (tail about 2 in length.," The object is trailed in both but, in the blue plate, shows a prominent diffuse tail about $\arcmin$ in length."
 The red plate shows only a hint of this tail., The red plate shows only a hint of this tail.
 The color (B-R = -]) is inconsistent with scattering [rom dust. but. suggests instead resonance fInorescent scattering [rom an ion tail.," The color (B-R = -1) is inconsistent with scattering from dust, but suggests instead resonance fluorescent scattering from an ion tail."
 The position angle of the tail. being about 15° [rom radial to the sun. is also more consistent with the expected direction of a plasma tail than with a dust tail.," The position angle of the tail, being about $\degr$ from radial to the sun, is also more consistent with the expected direction of a plasma tail than with a dust tail."
 Curiously. LOTP was re-observed on November 22 and 25 but then showed no trace of a dail (Cunningham 1950). ancl no comet-like activity has been reported since (Lowry ancl Weissman 2003. Ishiguro οἱ al.," Curiously, 107P was re-observed on November 22 and 25 but then showed no trace of a tail (Cunningham 1950), and no comet-like activity has been reported since (Lowry and Weissman 2003, Ishiguro et al."
 2011)., 2011).
 Dased on a statistical dvnamical model. Bottke et al. (," Based on a statistical dynamical model, Bottke et al. ("
2002) concluded that there is a chance that 107P is a captured Jupiter family comet.,2002) concluded that there is a chance that 107P is a captured Jupiter family comet.
 We consider a variety of processes capable of lanuching dust from a small body., We consider a variety of processes capable of launching dust from a small body.
 In each case. (he number of unknown but relevant physical parameters prevents anv exact treatment. but it remains instructive to consider the range of action of the mechanisms in the context of the asteroids.," In each case, the number of unknown but relevant physical parameters prevents any exact treatment, but it remains instructive to consider the range of action of the mechanisms in the context of the asteroids."
 since Whipple (1950). sublimation has been explored in great detail as the driver of mass loss from the classical comets.," Since Whipple (1950), sublimation has been explored in great detail as the driver of mass loss from the classical comets."
 It. need. not be re-described in detail here., It need not be re-described in detail here.
 Although simple in concept. detailed studies of comets show (hat sublimation is a remarkably complex process when [actors such as the porosity of the surface. nucleus rotation. the conduction of heat into (the interior and the development of a refractory mantle are considered et al.," Although simple in concept, detailed studies of comets show that sublimation is a remarkably complex process when factors such as the porosity of the surface, nucleus rotation, the conduction of heat into the interior and the development of a refractory mantle are considered (Guilbert-Lepoutre et al."
 2011)., 2011).
 One simplification possible for the present objects is (he assumption that asteroids contain no amorphous ice. since temperatures in the asterokl belt are too high for it to escape ervstallization.," One simplification possible for the present objects is the assumption that asteroids contain no amorphous ice, since temperatures in the asteroid belt are too high for it to escape crystallization."
 Aecordinely. we address only the highly idealized case of sublimation from an exposed crystalline ice surface in thermal equilibrium wilh sunlight.," Accordingly, we address only the highly idealized case of sublimation from an exposed crystalline ice surface in thermal equilibrium with sunlight."
 The sublimation mass flix per unit area. din/d! (kg ms 1) from a patch of surface," The sublimation mass flux per unit area, $dm/dt$ (kg $^{-2}$ $^{-1}$ ) from a patch of surface"
It is convenient to express the density and radius m terms of the dimensionless Launc-Eiidoeu functions 0 aud & where e has units of leneth aud is defined by (sec. e.g... Shapiro Teukolsky 1983).,"It is convenient to express the density and radius in terms of the dimensionless Lane-Emden functions $\theta$ and $\xi$ where $a$ has units of length and is defined by (see, e.g., Shapiro Teukolsky 1983)."
 In terius of these quautities. the mass ar) can be written where the prime denotes a derivative with respect to £.," In terms of these quantities, the mass $m(r)$ can be written where the prime denotes a derivative with respect to $\xi$."
 For --3. the values of© aud 0’ at the surface of the star are £1=6.897 and |0/(£4)|=0.012.," For $n=3$, the values of $\xi$ and $\theta'$ at the surface of the star are $\xi_1 = 6.897$ and $|\theta'(\xi_1)| = 0.0424$."
 Using (77)) and (79)). we find that (76)) reduces to a criterion on 0(£) aloue: where we have adopted (RIAMair~150 for the critical configuration.," Using \ref{LE_var}) ) and \ref{LE_mass}) ), we find that \ref{crit1_1}) ) reduces to a criterion on $\theta'(\xi)$ alone: where we have adopted $(R/M)_{\rm crit} \sim 450$ for the critical configuration."
 We plot 0'(£)/0'(£4) as a function of€ iu the top pancl of Fieure 6. and mark the threshold value 2.887 bw the horizoutal ine.," We plot $\theta'(\xi)/\theta'(\xi_1)$ as a function of $\xi$ in the top panel of Figure 6, and mark the threshold value $2.887$ by the horizontal line."
 A particle at a radius for which 0'(£)/0'(£4) is ercater than this value cau be captured. if the regious interior to this radius have collapsed to a black hole.," A particle at a radius for which $\theta'(\xi)/\theta'(\xi_1)$ is greater than this value can be captured, if the regions interior to this radius have collapsed to a black hole."
 For €2Gua~OST. O'(£3/0'(£4) is less than the capture hreshold. aud therefore the augular momentum barrier would prevent these regions from beiug caught by a newly ornmed. mterior black hole.," For $\xi > \xi_{\rm max} \sim 0.57 \,\xi_1$, $\theta'(\xi)/\theta'(\xi_1)$ is less than the capture threshold, and therefore the angular momentum barrier would prevent these regions from being caught by a newly formed, interior black hole."
 For the outermost regions. lis is not surprising. since the coufiguratiou is critically rotating.," For the outermost regions, this is not surprising, since the configuration is critically rotating."
 Therefore. the outer region may remain iu orbit. xhaps in a ciretuustellar disk. even if the rest of the star was collapsed.," Therefore, the outer region may remain in orbit, perhaps in a circumstellar disk, even if the rest of the star has collapsed."
 Iowever. frou the middle panel in Figure 6 we find that this outer region ouly coutaius of the nass. while about of the mass could form a black role.," However, from the middle panel in Figure 6 we find that this outer region only contains $~5\%$ of the mass, while about of the mass could form a black hole."
 This. again. is a consequence of 5j—3 polvtropes iue extremely ceutrally coudeused.," This, again, is a consequence of $n=3$ polytropes being extremely centrally condensed."
 According to Fieure 6. the aneular momentum barricr would also prevent particles at very all radii ©<0.05. youn being captured by a black hole interior to that radius.," According to Figure 6, the angular momentum barrier would also prevent particles at very small radii, $\xi \lesssim 0.05 \xi_1$, from being captured by a black hole interior to that radius."
 However. for this to be relevaut. the initial black hole would have to be restricted to a very suiall fractional size. and as we will see below. the secoud criterion (73)) docs uot allow such siunall black holes to form.," However, for this to be relevant, the initial black hole would have to be restricted to a very small fractional size, and as we will see below, the second criterion \ref{crit2}) ) does not allow such small black holes to form."
 We may reverse the above argent aud view eq. (80)), We may reverse the above argument and view eq. \ref{crit1_2}) )
 as a condition. on Z/M., as a condition on $R/M$.
 Since the left haud side. (0(£)/0'(£4). has a maxim of about 6.6. R/AL las to ο smaller than about 2350 for black bole formation o be possible.," Since the left hand side, $\theta'(\xi)/\theta'(\xi_1)$, has a maximum of about 6.6, $R/M$ has to be smaller than about 2350 for black hole formation to be possible."
 It is interesting that R/A of the critical configuration <A is only about five times smaller hau this threshold compaction. just barely allowing the supermassive star to form a supermassive black hole.," It is interesting that $R/M$ of the critical configuration $A$ is only about five times smaller than this threshold compaction, just barely allowing the supermassive star to form a supermassive black hole."
 This argunient sugeests that primordial eas mayhave to pass hrough a phase as a SMS where it can lose angular uonientun before it can possibly collapse to a SMDIT., This argument suggests that primordial gas mayhave to pass through a phase as a SMS where it can lose angular momentum before it can possibly collapse to a SMBH.
 We cun similarly evaluate the second criterion. eq. (73)).," We can similarly evaluate the second criterion, eq. \ref{crit2}) ),"
 in terms of Lauc-Euidenu functions., in terms of Lane-Emden functions.
 The angular nonmentum Jtr) of the matter euclosed within radius r.," The angular momentum $J(r)$ of the matter enclosed within radius $r$ ,"
of model spectra provided. ancl the bolometric corrections (DC) are taken trom (2000).,"of model spectra provided, and the bolometric corrections (BC) are taken from \cite{dril2000}."
. In (his table we also list (he upper limit of the stellar Iuminosity. assuming the solar Aj=4-4.74.," In this table we also list the upper limit of the stellar luminosity, assuming the solar $M_{\rm bol} = +4.74$."
" Similarly we can use the Oyoq limiting magnitude. and through calculate the colour dillerence eyµη as a Iunction of spectral (wpe. ancl hence calculate an equivalent Vj, and upper limit to the stellar luminosity."," Similarly we can use the $U_{\rm 300}$ limiting magnitude, and through calculate the colour difference $c_{\rm V-300}$ as a function of spectral type, and hence calculate an equivalent $M_{\rm bol}$ and upper limit to the stellar luminosity."
 The limits we have set on the total luminosity of PSN1999e1 allow comparison will stellar evolutionary model predictions Lor pre-supernova massive stars. aud we have chosen the Z=0.04 metallicity Geneva tracks (Mevnetetal.1994:Schaller1992.withtwiceihenormalmass-lossrates) for the lollowing reason.," The limits we have set on the total luminosity of PSN1999gi allow comparison with stellar evolutionary model predictions for pre-supernova massive stars, and we have chosen the Z=0.04 metallicity Geneva tracks \citep[with twice the normal mass-loss rates]{mey94,sch92}
 for the following reason."
" The oxvgen abundance gradient in NGC3184 has been determined by Zaritskyetal.(1994). and their region 68""N and Q""E (2.3kkpe radial distance) Irom (he nucleus of NGC3184 is almost certainly the reeion associated wilh NGC3184-ODI.", The oxygen abundance gradient in NGC3184 has been determined by \cite{zar94} and their region $68''$ N and $0''$ E kpc radial distance) from the nucleus of NGC3184 is almost certainly the region associated with NGC3184-OB1.
 The Ο/Η abundance derived is 9.2626 ddex. suggesting that the stars in this cluster are significantly. more metal rich. than solar (3.83dex[romSauval 1993).," The O/H abundance derived is $\pm$ dex, suggesting that the stars in this cluster are significantly more metal rich than solar \citep[8.83\,dex from][]{grev98}."
. The mass-loss rates used in these models are particularly hieh for the more massive stars which experience Woll-Havet. phases of evolution., The mass-loss rates used in these models are particularly high for the more massive stars which experience Wolf-Rayet phases of evolution.
 ILowever for masses of less than LOA. (which we show below is the region of interest) mass-loss plavs only a minor role in determining the point in the II-R. diagram at which the star explocles and we find similar constraints wilh alternative models. including those with no mass loss (e.g.Polsetal.1993).," However for masses of less than $10\,M_\odot$ (which we show below is the region of interest) mass-loss plays only a minor role in determining the point in the H-R diagram at which the star explodes and we find similar constraints with alternative models, including those with no mass loss \citep[e.g.][]{pols98}."
. The ΑΕΡΟΣ pre-explosion images are sensitive to all objects located in the shaded regions in 33., The WFPC2 pre-explosion images are sensitive to all objects located in the shaded regions in 3.
 For reference we show the position of Sk—69202. the B3la progenitor of SN1987À. Clearly PSN1999e1 was not a similar massive D-tvpe progenitor. or Π would have been detected on both the FGOGW and F300W frames.," For reference we show the position of $-$ 69202, the B3Ia progenitor of SN1987A. Clearly PSN1999gi was not a similar massive B-type progenitor, or it would have been detected on both the F606W and F300W frames."
 The best. and fairly conservative. estimate of the upper mass of the progenitor is 5M...," The best, and fairly conservative, estimate of the upper mass of the progenitor is $^{+3}_{-2}$ $_{\odot}$ ."
 We have assumed a, We have assumed a
Effects of disk accretion on structure of vouug stars have been investigated by Alercer-Suuith (19851). Palla Stabler (1992). Siess Forestini (1996). IIartiuaun (1097). Siess (1997. 1999).,"Effects of disk accretion on structure of young stars have been investigated by Mercer-Smith (1984), Palla Stahler (1992), Siess Forestini (1996), Hartmann (1997), Siess (1997, 1999)."
 Some of these authors studies how the leat advected into the star with the freshly accreted material iffects protostelhu properties., Some of these authors studies how the heat advected into the star with the freshly accreted material affects protostellar properties.
 However. noue of these iuvestigations looked at the effect of heat deposited i the stellar bv radiation originating iu the Inner parts of the circumstellar disk. where most of the accretion cucrey is released (see Figure P. for a schematic representation).," However, none of these investigations looked at the effect of heat deposited at the stellar by radiation originating in the inner parts of the circumstellar disk, where most of the accretion energy is released (see Figure \ref{fig:scheme} for a schematic representation)."
 Caven that accretion bhuninositv may easily exceed the iutrinsic stellar Iuuinositv (buuinositv derived from gravitational contraction. cooling aud. possibly. deuterimnu burning in stellar iuterior). onission of this effect nav not be justified in nisu cases.," Given that accretion luminosity may easily exceed the intrinsic stellar luminosity (luminosity derived from gravitational contraction, cooling and, possibly, deuterium burning in stellar interior), omission of this effect may not be justified in many cases."
 Iu this paper we investigate stellar irradiation bv the civetuustellar disk and address the importance of this effect in deterimuniug the intrinsic hunuinositv of voung stars., In this paper we investigate stellar irradiation by the circumstellar disk and address the importance of this effect in determining the intrinsic luminosity of young stars.
 We calculate the spatial distribution of the disk flux on the stellay surface aud determine when irradiation is inportaut iu 82.., We calculate the spatial distribution of the disk flux on the stellar surface and determine when irradiation is important in \ref{sect:T_dist}.
 The effect of nradiation on stellar cooling is investigated locally in 83. aud elobally in &L.., The effect of irradiation on stellar cooling is investigated locally in \ref{sect:stellar_cool} and globally in \ref{sect:total_cooling}.
 Finally. in $5. we discuss the applications of this study to some real objects and its possible liuitatious.," Finally, in \ref{sect:disc} we discuss the applications of this study to some real objects and its possible limitations."
" We start by calculating the distribution on the stellar surface of the radiative fux. £j, produced by the disk.", We start by calculating the distribution on the stellar surface of the radiative flux $F_{irr}$ produced by the disk.
" We consider an axisvuuuetric ecometrically thin disk accreting onto a star with radius A, aud mass M,.", We consider an axisymmetric geometrically thin disk accreting onto a star with radius $R_\star$ and mass $M_\star$.
 Flux Fi dutercepted by the star is a function of 0. tho angle between the normal to the stellar surface aud the normal to the disk (coincident with the polar axis of the star. assunune that disk lies im the stellar equatorial plac).," Flux $F_{irr}$ intercepted by the star is a function of $\theta$ – the angle between the normal to the stellar surface and the normal to the disk (coincident with the polar axis of the star, assuming that disk lies in the stellar equatorial plane)."
 Polar regious of the star are exposed to the radiation of ouly the distant. cool parts of the disk. while the equatorial regions have a direct view to the iunerimost parts of the disk where most of the energy is dissipated.," Polar regions of the star are exposed to the radiation of only the distant, cool parts of the disk, while the equatorial regions have a direct view to the innermost parts of the disk where most of the energy is dissipated."
" One can easily show that a disk extending all the wav to the stellar surface gives rise to irradiation flux £5,(0) eiven by (Adis Shu 1986: Popliun 1997) ΕΠ] where Ris the evlindrical radius. cos=δνπι]. RR,=RfcosO. aud F(R) is the energyo. raciated by the unit surface area of the disk per unit of time."," One can easily show that a disk extending all the way to the stellar surface gives rise to irradiation flux $F_{irr}(\theta)$ given by (Adams Shu 1986; Popham 1997) F_d(R)R where $R$ is the cylindrical radius, $\cos\phi_c=R_\star/(r\sin\theta)$, $R_{in}=R_\star/\cos\theta$, and $F_d(R)$ is the energy radiated by the unit surface area of the disk per unit of time."
 In Appeudix A we demonstrate that this expression can be reduced to a one-dimensional integral which is easier to analyze than equation (01) , In Appendix A we demonstrate that this expression can be reduced to a one-dimensional integral which is easier to analyze than equation \ref{eq:irr_flux}) ).
To find the explicit dependence of ἔτι ou 0 one needs to know Fiy(2) which is determined by the viscous dissipation in the disk., To find the explicit dependence of $F_{irr}$ on $\theta$ one needs to know $F_d(R)$ which is determined by the viscous dissipation in the disk.
 Studies of steady-state thin accretion disks have ecnerally found that where AZ is a iuass accretion rate and the function F(R). enibodyiug the details of the disk emissivitv near the stellar surface. behaves as fo>1 when Π R..," Studies of steady-state thin accretion disks have generally found that where $\dot M$ is a mass accretion rate and the function $f(R)$, embodying the details of the disk emissivity near the stellar surface, behaves as $f\to 1$ when $R\gg R_\star$."
 With Fy given by (6)) one finds where the dineusiouless function g(7) is given by equation (AT)).," With $F_d$ given by \ref{eq:vis_dissip}) ) one finds ), where the dimensionless function $g(\theta)$ is given by equation \ref{eq:g}) )."
" A standard disk with zero torque at the stellar surface (situation appropriate for accretion onto black holes) has (Shakura Suuvaev 1973) f(R)=1.(RL/R)V7,", A standard disk with zero torque at the stellar surface (situation appropriate for accretion onto black holes) has (Shakura Sunyaev 1973) $f(R)=1-(R_\star/R)^{1/2}$.
" The total viscous dissipation in such a disk is Ly=(L/2GALAL/R, aud the eas at the iuner edec of the disk rotates at the local Isepleriau velocity.", The total viscous dissipation in such a disk is $\dot E_d = (1/2)GM_\star \dot M/R_\star$ and the gas at the inner edge of the disk rotates at the local Keplerian velocity.
" This is Inappropriate in our case since the eas speed has to match the velocity of the stellar surface at R=R, (for simplicity assumed to be zero in our case).", This is inappropriate in our case since the gas speed has to match the velocity of the stellar surface at $R=R_\star$ (for simplicity assumed to be zero in our case).
 As a result a boundary laver must form uear the stellar surface iu which the azimuthal velocity of the gas is lowered bv the viscous torque from the local I&epleriau value to the stellar rotation speed., As a result a boundary layer must form near the stellar surface in which the azimuthal velocity of the gas is lowered by the viscous torque from the local Keplerian value to the stellar rotation speed.
 Viscous dissipation dramatically increases. gas temperature in this laver creating an additional source of radiative flux very close to the stellar surface., Viscous dissipation dramatically increases gas temperature in this layer creating an additional source of radiative flux very close to the stellar surface.
 Irracliation bv the boundary ατα ciission boosts up the stellar surface temperature dmi a narrow, Irradiation by the boundary layer emission boosts up the stellar surface temperature in a narrow
spectroscopy. (e.g.Petitetal.2003)...,spectroscopy \citep[e.g.][]{2008MNRAS.388...80P}.
 A comparison of these observations aud numerical caleulations of the stellar diamo could give new insight into the stellar magnetic field., A comparison of these observations and numerical calculations of the stellar dynamo could give new insight into the stellar magnetic field.
 Finally. our stellar: ANID dynamo study would also contribute to the understanding of recent investigations into stellar magnetic evelic activity periods Brandenburg 1999)..," Finally, our stellar MHD dynamo study would also contribute to the understanding of recent investigations into stellar magnetic cyclic activity periods \citep{1984ApJ...287..769N,1999ApJ...524..295S}."
 We are most grateful to Dr. M. Rempel for helpful advice., We are most grateful to Dr. M. Rempel for helpful advice.
 Numerical computations were carried out at the General-Purpose PC farm in the Center for Computational Astrophysics (CICA) of the National Astronomical Observatory of Japan., Numerical computations were carried out at the General-Purpose PC farm in the Center for Computational Astrophysics (CfCA) of the National Astronomical Observatory of Japan.
 The page charge for this paper is supported by CICA., The page charge for this paper is supported by CfCA.
 We have greatly benefited from the proofreading/editing assistance from the GCOE program., We have greatly benefited from the proofreading/editing assistance from the GCOE program.
"7 and 7 measure the projected correlation ΠΙΟΙΟ wy(r,) for this sample on small and intermeciate scales: We follow ? and set πω.=80 Mpce/h. which is large enough to include most correlated pairs and. produce stable estimates of ie,(7,)).","\citet{zehavi/etal:2005a} and \citet{masjedi/etal:2006} measure the projected correlation function $w_p(r_p)$ for this sample on small and intermediate scales: We follow \citet{zehavi/etal:2005a} and set $\pi_{max} = 80$ $h$, which is large enough to include most correlated pairs and produce stable estimates of $w_p(r_p)$."
" 2 recover missing Liber collision pairs by computing wy(r,) by cross correlation between the SDSS spectroscopic and imaging samples.", \citet{masjedi/etal:2006} recover missing fiber collision pairs by computing $w_p(r_p)$ by cross correlation between the SDSS spectroscopic and imaging samples.
" Dhey also correct for photometric biases of close galaxy pairs. which can introduce incompleteness of pairs with separation ry,S0.1 Mpc/h. We present the projected correlation function τρ) averaged over 20 mock catalogs produced wilh our SO halo catalog using our maximum likelihood ILOD in Figure 5.."," They also correct for photometric biases of close galaxy pairs, which can introduce incompleteness of pairs with separation $r_p \lesssim 0.1$ $h$ We present the projected correlation function $w_p(r_p)$ averaged over 20 mock catalogs produced with our SO halo catalog using our maximum likelihood HOD in Figure \ref{fig:wprp1}."
 We lind excellent agreement wilh the measurements of ?.., We find excellent agreement with the measurements of \citet{masjedi/etal:2006}.
 Using the diagonal error bars reported in ?.. we find 4?=7.5 for the outer 15 points.," Using the diagonal error bars reported in \citet{masjedi/etal:2006}, we find $\chi^2 = 7.5$ for the outer 15 points."
" There is substantial discrepancy with the inner 3 points al rj,=0.01.0.016.0.026 (not shown in Fig. 3)):"," There is substantial discrepancy with the inner 3 points at $r_p = 0.01, 0.016, 0.026$ (not shown in Fig. \ref{fig:wprp1}) );"
 P=29 for all 18 points., $\chi^2 = 29$ for all 18 points.
 The discrepancy is not surprising since (hese small distances are comparable to the force resolution of our simulation., The discrepancy is not surprising since these small distances are comparable to the force resolution of our simulation.
" Though the CiC method relies primarily on pairs with r,X0.8 Mpc/h. our mock catalogs reproduce the features of the observed ορ) by adjusting a single parameter 2;5,3 to match the large scale (o20 Mpc/h) bias probed by wy)(rp)."," Though the CiC method relies primarily on pairs with $r_p \leq 0.8$ $h$, our mock catalogs reproduce the features of the observed $w_p(r_p)$ by adjusting a single parameter $\sigma_{log M}$ to match the large scale $\sim 20$ $h$ ) bias probed by $w_p(r_p)$."
" Note that a sharp (ransilion [rom OQ to 1 for 4,7) can be ruled out with confidence.", Note that a sharp transition from 0 to 1 for $N_{cen}(M)$ can be ruled out with confidence.
" Figure 8. shows p(y) for catalogs with Goya,=0.2. 0.7. and 1.3 for comparison."," Figure \ref{fig:wprp1} shows $w_p(r_p)$ for catalogs with $\sigma_{log M} = 0.2$, 0.7, and 1.3 for comparison."
" All three catalogs match the observed clustering at o,Z50.5 Mpc//h where we have CiC constraints. but only catalogs with σον0.1 match the observed clustering on e2—20 Mpc/h scales. the regime where two-halo pairs"," All three catalogs match the observed clustering at $r_p \lesssim 0.8$ $h$ where we have CiC constraints, but only catalogs with $\sigma_{log M} \sim 0.7$ match the observed clustering on $\sim 2-20$ $h$ scales, the regime where two-halo pairs"
the odf files with the SAS software (version 7.0.0).,the odf files with the SAS software (version 7.0.0).
 Given its higher sensitivity. we use the time average EPIC/pn spectrum for the analysis of each object. except for IGR J16482-3036 for which only EPIC/MOS data are available.," Given its higher sensitivity, we use the time average EPIC/pn spectrum for the analysis of each object, except for IGR J16482-3036 for which only EPIC/MOS data are available."
 X-ray events corresponding to patterns 0-12 and 0-3 were selected from the MOS and pn. respectively.," X-ray events corresponding to patterns 0-12 and 0-4 were selected from the MOS and pn, respectively."
 We used the most updated calibration files available at the time of the reduction for each source data., We used the most updated calibration files available at the time of the reduction for each source data.
 Source light curves and spectra were extracted from circular regions of typically 50” centered on the source. while background products were obtained from off-set regions close to the source.," Source light curves and spectra were extracted from circular regions of typically $\arcsec$ centered on the source, while background products were obtained from off-set regions close to the source."
 Exposures have been filtered for periods of high background and the effective exposures are reported in Table 2 as well as the observation date. the pn filter and the number of counts per bin used to rebin the spectral channels.," Exposures have been filtered for periods of high background and the effective exposures are reported in Table \ref{table=obs_info} as well as the observation date, the pn filter and the number of counts per bin used to rebin the spectral channels."
 Spectra were binned according to the luminosity of each source., Spectra were binned according to the luminosity of each source.
 The ancillary and detector response matrices were generated using the SAS and. tasks., The ancillary and detector response matrices were generated using the SAS and tasks.
 The and data were fitted together and analyzed using XSPEC v.12.4.0., The and data were fitted together and analyzed using XSPEC v.12.4.0.
 Since the and observations are not simultaneous. à cross-calibration constant C has been introduced 1n our best-fit models.," Since the and observations are not simultaneous, a cross-calibration constant $C$ has been introduced in our best–fit models."
 This has been done to take into account possible cross-calibration mismatches between the two instruments or variability in the sources., This has been done to take into account possible cross-calibration mismatches between the two instruments or variability in the sources.
 The constant was left free to vary and. for each fit. its value is reported in the relevant Table.," The constant was left free to vary and, for each fit, its value is reported in the relevant Table."
 Galactic absorption is implicitly included in all spectral models: abundances are those of Anders Grevesse (1989)., Galactic absorption is implicitly included in all spectral models; abundances are those of Anders Grevesse (1989).
 The errors. lower and upper limits quoted correspond to confidence range for one interesting parameter (ie. Ay= 2.7]: Avni 1976).," The errors, lower and upper limits quoted correspond to confidence range for one interesting parameter (i.e. $\Delta\chi^2 = 2.71$ ; Avni 1976)."
 The broad-band 0.5-150 keV and spectrum. of each source has been initially fitted with a power-law model absorbed by intrinsic cold absorption. plus a soft X-ray component and à narrow Gaussian emission (Fe) line.," The broad-band 0.5-150 keV and spectrum of each source has been initially fitted with a power–law model absorbed by intrinsic cold absorption, plus a soft X-ray component and a narrow Gaussian emission (Fe) line."
" A simple parameterization has been employed to model the soft component found to be present in six of our sources: either a black body or MEKAL thermal plasma model with temperature AT provided a good description of the data. except for LEDA 168563. where a soft power-law model (L,,;5,) was instead preferred."," A simple parameterization has been employed to model the soft component found to be present in six of our sources: either a black body or a MEKAL thermal plasma model with temperature $kT$ provided a good description of the data, except for LEDA 168563, where a soft power-law model $\Gamma_{soft}$ ) was instead preferred."
 All detected FeK« emission lines were found to be consistent with à narrow Gaussian profile. so that the line width was fixed to c = 10 eV. In a few sources. the quality of the fit improves significantly with the introduction of additional spectral components such as a partial covering absorption model in XSPEC) in 4U 1344-60 and IGR 116558-5203.," All detected $\alpha$ emission lines were found to be consistent with a narrow Gaussian profile, so that the line width was fixed to $\sigma$ $=$ 10 eV. In a few sources, the quality of the fit improves significantly with the introduction of additional spectral components such as a partial covering absorption model in XSPEC) in 4U 1344-60 and IGR J16558-5203."
 An extra Gaussian emission line is instead required in IGR. J17418-1212 and IGR J18027-1455: finally an absorption edge was required in the case of FRL 1146 (AO4)., An extra Gaussian emission line is instead required in IGR J17418-1212 and IGR J18027-1455; finally an absorption edge was required in the case of FRL 1146 (AO4).
 In Table 3 we show the best-fit parameters together with the model fluxes in the 2-10 keV and 20-100 keV bands., In Table \ref{table=po} we show the best–fit parameters together with the model fluxes in the 2-10 keV and 20-100 keV bands.
" Although the best-fit solutions are obtained by including the additional spectral components described above. the values of their parameters are not included in Table 3 but are discussed in detail. for each source. in the next subsection,"," Although the best–fit solutions are obtained by including the additional spectral components described above, the values of their parameters are not included in Table \ref{table=po} but are discussed in detail, for each source, in the next subsection."
 To check for the presence of a high energy cut-off. we replaced the simple power-law in the best-fit model used in Table 3.. with one having an exponential cut-off(cutoffp/ model in XSPEC) and report in Table 4. the results provided by this change on the main fit parameters.," To check for the presence of a high energy cut-off, we replaced the simple power-law in the best-fit model used in Table \ref{table=po}, with one having an exponential cut-off model in XSPEC) and report in Table \ref{table=cut} the results provided by this change on the main fit parameters."
 Finally. to also take into account the possible presence of a reflection component. we replaced the exponentially cut-off power-law with the model in XSPEC.," Finally, to also take into account the possible presence of a reflection component, we replaced the exponentially cut-off power-law with the model in XSPEC."
 The reflection component is described by two fundamental parameters:Q/27. which ts the fraction of a neutral. plane parallel slab illuminated by the power-law photons. and or the slab inclination angle with respect to the line of sight.," The reflection component is described by two fundamental parameters:, which is the fraction of a neutral, plane parallel slab illuminated by the power-law photons, and or the slab inclination angle with respect to the line of sight."
 Since all parameters turned out to be identical to those obtained with left free to vary. we choose to freeze to 0.45.," Since all parameters turned out to be identical to those obtained with left free to vary, we choose to freeze to 0.45."
 In Table 5 we report the values obtained for the main spectral parameters., In Table \ref{table=pex} we report the values obtained for the main spectral parameters.
 In the next subsection. we briefly describe. for each individual source. the best-fit models and the values of the additional spectral components which. for the sake of clarity.," In the next subsection, we briefly describe, for each individual source, the best–fit models and the values of the additional spectral components which, for the sake of clarity,"
The center of AIST was imaged with the ΛΕΡΟΣ as part of a 30 orbit program in May and June of 2001.,The center of M87 was imaged with the WFPC2 as part of a 30 orbit program \citep{Baltz} in May and June of 2001.
 During each orbit. [our exposures totalling 10405 were taken in the FsltW filter. with a single matching 400s exposure in the FGOGW filler.," During each orbit, four exposures totalling 1040s were taken in the F814W filter, with a single matching 400s exposure in the F606W filter."
 The four FS14W images were dithered by half pixel steps to allow for the images to be interlaced directlv into a 2x2 grid., The four F814W images were dithered by half pixel steps to allow for the images to be interlaced directly into a 2x2 grid.
 Average images were created from the final full data set by stacking ihe images. alter they. had. been sinc-interpolated to a common origin.," Average images were created from the final full data set by stacking the images, after they had been sinc-interpolated to a common origin."
 The images [rom dav 12 were excluded. from (his step as they. had more jitter (han the other images., The images from day 12 were excluded from this step as they had more jitter than the other images.
 This combination creates images (hat have total exposure (times of 301608 for FES14W. and 110005 [or F606W. Cluster detection was performed separately on the images for each filler using Source Extractor (Bertin&Arnouts1996)..," This combination creates images that have total exposure times of $30160$ s for F814W, and $11600$ s for F606W. Cluster detection was performed separately on the images for each filter using Source Extractor \citep{sextractor}."
. As we are imaging the center of the galaxy. the galaxy profile is steep.," As we are imaging the center of the galaxy, the galaxy profile is steep."
 Because the main source of noise in the image is from the variation in the nunmber of photons detected from the galaxy light. (his steep profile creates a rapid change in (he noise level across the images.," Because the main source of noise in the image is from the variation in the number of photons detected from the galaxy light, this steep profile creates a rapid change in the noise level across the images."
 To prevent the change in noise from altering how clusters are detected at different distances [rom the galaxy. we created an image to model the background level of the galaxy itself.," To prevent the change in noise from altering how clusters are detected at different distances from the galaxy, we created an image to model the background level of the galaxy itself."
 This was done by median filtering the image with a box larger than the size of anv of the globular cluster features., This was done by median filtering the image with a box larger than the size of any of the globular cluster features.
 As the radii of the largest clusters can be up (o 20 pixels. we use a filter box of 50 x 50 pixels.," As the radii of the largest clusters can be up to 20 pixels, we use a filter box of 50 x 50 pixels."
 This background image was used (o weight the detection threshold. which helps minimize radial variations in our object detection.," This background image was used to weight the detection threshold, which helps minimize radial variations in our object detection."
 since the faint end of the observed GCLF is sensitive to the contamination of objects that are not globular clusters. we want (o ensure that we remove as many of these objects as possible.," Since the faint end of the observed GCLF is sensitive to the contamination of objects that are not globular clusters, we want to ensure that we remove as many of these objects as possible."
 The majority of the very. faintest objects detected are likely to be simply peaks in the variation of the background., The majority of the very faintest objects detected are likely to be simply peaks in the variation of the background.
 As the background. subtracted image histograms are svimneltric about zero. we expect that the magnitude and frequency of these variations are the same between peaks and valleys.," As the background subtracted image histograms are symmetric about zero, we expect that the magnitude and frequency of these variations are the same between peaks and valleys."
 TDherelore. by inverting the sign of the data images (simplv multiplving by —1). and searching on (hese images. we can determine how manv candidate clusters are caused by the variations in the background. level.," Therefore, by inverting the sign of the data images (simply multiplying by $-1$ ), and searching on these images, we can determine how many candidate clusters are caused by the variations in the background level."
 Based on these fests. we chose detection thresholds of 30 with a minimum object size of 2 pixels.," Based on these tests, we chose detection thresholds of $3\sigma$ with a minimum object size of 2 pixels."
 Although ilis tempting to push this lower to increase (he completion al fainter magnitudes. below this limit the increase in (he umber of objects detected on the data image is closely matched by," Although it is tempting to push this lower to increase the completion at fainter magnitudes, below this limit the increase in the number of objects detected on the data image is closely matched by"
"The diagnostic unknowns ©, and i, are given by and With the notation 9=(ay.vo). the nonlinear equations of the zero mode can be written as and V and div are the 2D gradient and divergence operators. respectively.","The diagnostic unknowns $\phi_n$ and $w_n$ are given by and With the notation $\vb=(u_0,\,v_0)^T$, the nonlinear equations of the zero mode can be written as Here and $\triangledown$ and $\udiv$ are the 2D gradient and divergence operators, respectively."
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2.," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2.9," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2.9)," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2.9).," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
 Without considering other modes. the nonlinear equations ofthe zero mode become where ο—ὦμ+[UVHy. We supplement the svstem with the following boundary conditions: The wellposedness of the linearized svstem associated with (2.9)..," Without considering other modes, the nonlinear equations ofthe zero mode become where $\varphi = \phi_0 + f\bar{U}_0 \sqrt {H} y$ We supplement the system with the following boundary conditions: The well–posedness of the linearized system associated with ,"
that the errors on Zeeman splitting determination in VOS should be increased.,that the errors on Zeeman splitting determination in V08 should be increased.
 For the majority of the masers in. VO8. this increase is less than a factor of ~5.," For the majority of the masers in V08, this increase is less than a factor of $\sim5$."
 However. for the masers of 0ο620.390 the errors need to be increased with a factor of 40 and 8 for the main and secondary maser features respectively.," However, for the masers of G09.62+0.20 the errors need to be increased with a factor of $40$ and $8$ for the main and secondary maser features respectively."
 We determined the Zeeman splitting of the two brightest 6.7 GHz methanol maser features of GO9.62+0.20 using the RCP-LCP cross-correlation method as described in VO8., We determined the Zeeman splitting of the two brightest 6.7 GHz methanol maser features of G09.62+0.20 using the RCP-LCP cross-correlation method as described in V08.
 Fig., Fig.
 3 shows the line-of-sight magnetic field strength Bj for the two maser features at each epoch as derived from the Zeeman splitting AV. using 0.049 G! as the best estimate for the Zeeman splitting coefficient.," \ref{Fig:var} shows the line-of-sight magnetic field strength $B_{\rm ||}$ for the two maser features at each epoch as derived from the Zeeman splitting $\Delta V_Z$, using $0.049$ $^{-1}$ as the best estimate for the Zeeman splitting coefficient."
 In addition to the monitoring epochs. the figure also includes the folded in observations presented in VOS.," In addition to the monitoring epochs, the figure also includes the folded in observations presented in V08."
 We find that the line-of-sight magnetic field strength in the secondary maser feature is stable during the observations. with an error weighted average magnetic field of Bj=10.9x2.3 mG. In contrast. the Zeeman splitting of the main maser feature sharply decreases at the time the maser flare reaches peak flux.," We find that the line-of-sight magnetic field strength in the secondary maser feature is stable during the observations, with an error weighted average magnetic field of $B_{\rm ||}=10.9\pm2.3$ mG. In contrast, the Zeeman splitting of the main maser feature sharply decreases at the time the maser flare reaches peak flux."
 Monitoring at weekly intervals has proven to be too coarse to put strong constraints on the length of the period with decreased AVz., Monitoring at weekly intervals has proven to be too coarse to put strong constraints on the length of the period with decreased $\Delta V_Z$.
 However. while the typical duration of a flare is up to two months. the decrease in Zeeman splitting lasts only for an approximate two week period around the peak of the flare.," However, while the typical duration of a flare is up to two months, the decrease in Zeeman splitting lasts only for an approximate two week period around the peak of the flare."
 Determining a error weighted average magnetic field for the main maser feature using the 7 epochs on either side of the two week period with decreased Zeeman splitting. we find Bj=11.0x2.2 mG. Thus. the magnetic field strength is remarkably similar on both masers. even though both features are separated by more than 200 AU etal..2005).," Determining a error weighted average magnetic field for the main maser feature using the 7 epochs on either side of the two week period with decreased Zeeman splitting, we find $B_{\rm
 ||}=11.0\pm2.2$ mG. Thus, the magnetic field strength is remarkably similar on both masers, even though both features are separated by more than 200 AU \citep{Goedhart05}."
., Fig.
 Fig. 5. thus indicates that. when the flare of the main maser feature reaches its peak flux. the Zeeman splitting decreases significantly and potentially even changes sign.," \ref{Fig:var} thus indicates that, when the flare of the main maser feature reaches its peak flux, the Zeeman splitting decreases significantly and potentially even changes sign."
" As the observations of VOS that also revealed a much lower (and possibly reversed) AV, were taken close to the peak of the previously flaring period. this behavior appears to repeat itself regularly."," As the observations of V08 that also revealed a much lower (and possibly reversed) $\Delta V_Z$ were taken close to the peak of the previously flaring period, this behavior appears to repeat itself regularly."
 No significant Zeeman splitting decrease 15 seen for the secondary maser feature., No significant Zeeman splitting decrease is seen for the secondary maser feature.
 The observations with the HartRAO telescope showed that the flare of this feature. while typically approximately 25 days after the flare of the main feature. was much more irregular and reached its peak at approximately the same time as the main feature.," The observations with the HartRAO telescope showed that the flare of this feature, while typically approximately 25 days after the flare of the main feature, was much more irregular and reached its peak at approximately the same time as the main feature."
" However. throughout the Effelsberg observations its flux variations were only ~10% while that of the primary feature was over 20%,"," However, throughout the Effelsberg observations its flux variations were only $\sim10\%$ while that of the primary feature was over $20\%$."
 The cause for the sudden decrease in Zeeman splitting during the maser flare is unknown., The cause for the sudden decrease in Zeeman splitting during the maser flare is unknown.
" We can confidently rule out instrumental effects as the reason for the decrease of AV, of the main maser feature for a number of reasons.", We can confidently rule out instrumental effects as the reason for the decrease of $\Delta V_Z$ of the main maser feature for a number of reasons.
 First of all. no significant corresponding decrease is found for the secondary naser feature at the same epochs.," First of all, no significant corresponding decrease is found for the secondary maser feature at the same epochs."
 Furthermore. the quite constant Zeeman splitting measured before and after the peak. and the fact that the negative Zeeman splitting was also found during the maser flare 8 months earlier. point to the stability of the instrumental setup as well as the robustness of the data reduction and analysis method.," Furthermore, the quite constant Zeeman splitting measured before and after the peak, and the fact that the negative Zeeman splitting was also found during the maser flare 8 months earlier, point to the stability of the instrumental setup as well as the robustness of the data reduction and analysis method."
 The observed effect is thus likely intrinsic to the source., The observed effect is thus likely intrinsic to the source.
" The measured decrease in AV, however. starts after the naser flux has already entered the flaring stage."," The measured decrease in $\Delta V_Z$ however, starts after the maser flux has already entered the flaring stage."
 Thus. while the measured Zeeman splitting variation is related to the naser flare. 1t is unlikely that the flare itself is caused by changes in the magnetic field.," Thus, while the measured Zeeman splitting variation is related to the maser flare, it is unlikely that the flare itself is caused by changes in the magnetic field."
 Here we deseribe three possible scenarios that could give rise to the observed effect., Here we describe three possible scenarios that could give rise to the observed effect.
 As. the observed Zeeman splitting is generated by the average magnetic field throughout, As the observed Zeeman splitting is generated by the average magnetic field throughout
Figure S shows the D and X-ray light curves normalised to their mean fluxes. where the B data has been shifted back by 6 days. according to the time lag measured.,"Figure \ref{norm_lcs} shows the B and X-ray light curves normalised to their mean fluxes, where the B data has been shifted back by 6 days, according to the time lag measured."
 Lt is clear that in this way the peaks in both light curves match very well. especially at. low optical Huxes.," It is clear that in this way the peaks in both light curves match very well, especially at low optical fluxes."
 The long term trend. has slightly higher amplitude in the B band while the amplitude of the rapid. Dluctuations is much larger in the X-rays. as expected from the analysis of the PDS.," The long term trend has slightly higher amplitude in the B band while the amplitude of the rapid fluctuations is much larger in the X-rays, as expected from the analysis of the PDS."
 The measured. lag of optical behind X-rav lluctuations suggests à reprocessing scenario where at least. part of the optical variability is produced by the variable X-ray heating., The measured lag of optical behind X-ray fluctuations suggests a reprocessing scenario where at least part of the optical variability is produced by the variable X-ray heating.
 In this case. the delay. interpreted as a light travel time. gives the distance between the X-ray. corona and the location of the reprocessor.," In this case, the delay, interpreted as a light travel time, gives the distance between the X-ray corona and the location of the reprocessor."
" The black hole mass of lis 3.10'M. (Petersonetal.2004). producing a ligh crossing time for l1 gravitational radius of LRfe=150 s. The time delay of optical behind N-ray luctuations. of G6c3 days therefore. corresponds to a light travel time through a distance of ~3500+1750R,.", The black hole mass of is $3\times 10^{7} M_\odot$ \citep{revmap} producing a light crossing time for 1 gravitational radius of $1 R_g/c=150$ s. The time delay of optical behind X-ray fluctuations of $\sim 6\pm 3$ days therefore corresponds to a light travel time through a distance of $\sim 3500\pm1750 R_g$.
 We note however that the region on the disc where the amount of reprocessec Hux peaks will not necessarily correspond to the region of peak intrinsic emission due to internal heating. since X-ray heating may dominate at larger radii as will be ciscusse below in Sec. 5.2.," We note however that the region on the disc where the amount of reprocessed flux peaks will not necessarily correspond to the region of peak intrinsic emission due to internal heating, since X-ray heating may dominate at larger radii as will be discussed below in Sec. \ref{intrinsicplusreprocessed}."
 Not all the optical variability can arise from reprocessed X-rays. however. as ds evident from the power spectrum. analvsis.," Not all the optical variability can arise from reprocessed X-rays, however, as is evident from the power spectrum analysis."
 The PDS of the DB. band. data shows that the variability. power continues to increase towards. lower frequencies. at. least down to frequencies 410? Lz. or correspondingly time-scales of ~300 days. while the Χ-pav variability power breaks and Uattens for time-scales below 2 cavs.," The PDS of the B band data shows that the variability power continues to increase towards lower frequencies, at least down to frequencies $4\times 10^{-8}$ Hz, or correspondingly time-scales of $\sim 300$ days, while the X-ray variability power breaks and flattens for time-scales below 2 days."
 ΙΓ all the optical variability was produced. by reprocessing. the X-ray. Hluctuations on time-scales shorter than 100 clays would have to be smoothed-out by dillerential light travel time to opposite sides of the reprocessor. requiring it to have a minimum raclius of ~50 light days.," If all the optical variability was produced by reprocessing, the X-ray fluctuations on time-scales shorter than 100 days would have to be smoothed-out by differential light travel time to opposite sides of the reprocessor, requiring it to have a minimum radius of $\sim50$ light days."
 In this case. the average delay. between optical and N-ravs would be of the same order. much longer than the measured 6 cay lag.," In this case, the average delay between optical and X-rays would be of the same order, much longer than the measured $\sim 6$ day lag."
 To exemplify this argument we constructed. transfer functions for dillerent reprocessing geometrics., To exemplify this argument we constructed transfer functions for different reprocessing geometries.
 I£ the optical variability. arises only from N-rav. reprocessing then the transfer function. has to produce both the smoothing of Iluctuations down to time-scales of ~300 davs and a delay of ~6 days between primary ancl reprocessecl Dux., If the optical variability arises only from X-ray reprocessing then the transfer function has to produce both the smoothing of fluctuations down to time-scales of $\sim 300$ days and a delay of $\sim 6$ days between primary and reprocessed flux.
 In this case. the variable part of the optical light curve would be olf)=.r(d)HD). where w(f) is the X-ray light curve. HE) ds the transfer. function. and denotes. convolution.," In this case, the variable part of the optical light curve would be $o(t)=x(t)\otimes i(t)$, where $x(t)$ is the X-ray light curve, $i(t)$ is the transfer function and $\otimes$ denotes convolution."
 Vhe relation of their Fourier transform would therefore be οι—NCP).HG). so that where JOLPF and INCA correspond. to the PDS of he optical and X-ray light curves. respectively. as caleulated in Sec. 3..," The relation of their Fourier transform would therefore be $O(f)=X(f)\times I(f)$, so that where $|O(f)|^2$ and $|X(f)|^2$ correspond to the PDS of the optical and X-ray light curves, respectively, as calculated in Sec. \ref{psd}."
 The plots in Figure 9 compare this PDS ratio (calculated. from the fits to the PDS using Ίσα., The plots in Figure \ref{transfer} compare this PDS ratio (calculated from the fits to the PDS using Eq.
 1 and shown here in green triple-dot-dashed. lines) to the |1Cf)E unctions predicted. for cach case., 1 and shown here in green triple-dot-dashed lines) to the $|I(f)|^2$ functions predicted for each case.
 We note that the PDS ratio was calculated for optical anc X-rav light) curves observed simultaneously. so the variations in the PDS shape expected due to the red noise nature of these light curves is not relevant.," We note that the PDS ratio was calculated for optical and X-ray light curves observed simultaneously, so the variations in the PDS shape expected due to the red noise nature of these light curves is not relevant."
 The reprocessor gcometries considered were a spherical shell distribution of small clouds. a flat optically thick disc and an optically thick disc which flares out at large raclii.," The reprocessor geometries considered were a spherical shell distribution of small clouds, a flat optically thick disc and an optically thick disc which flares out at large radii."
 These were chosen to mimic structures that probably exist in the vicinity of the black hole such as an optically thick accretion disc and the broad line region (BLIU) clouds., These were chosen to mimic structures that probably exist in the vicinity of the black hole such as an optically thick accretion disc and the broad line region (BLR) clouds.
 In all cases. model parameters were adjusted to produce the observed: average delav of 6 days between primary and reprocessed emission. calculated as the expectation value of the transfer function.," In all cases, model parameters were adjusted to produce the observed average delay of 6 days between primary and reprocessed emission, calculated as the expectation value of the transfer function."
 The spherical shell produces the strongest smoothing of long-term trends for a fixed time lag. compared to the other geometries.," The spherical shell produces the strongest smoothing of long-term trends for a fixed time lag, compared to the other geometries."
 This is because there is no delay between the intrinsic emission and. the response [rom the side of the sphere closest to the observer., This is because there is no delay between the intrinsic emission and the response from the side of the sphere closest to the observer.
 In disc geometries (except for an edge-on. view) the intrinsic Dux must first travel to the nearest side of the reprocessor. which is located out of the line of sight. delaving the start of the response.," In disc geometries (except for an edge-on view) the intrinsic flux must first travel to the nearest side of the reprocessor, which is located out of the line of sight, delaying the start of the response."
 Therefore. the same ‘smoothing power can be associated with a longer delay. if the disce is viewed face-on and the reprocessing region is concentrated in a narrow ring then there will be no smoothing of the response but the delay can be arbitrarily long. for an arbitrarily large ring radius.," Therefore, the same `smoothing power' can be associated with a longer delay, if the disc is viewed face-on and the reprocessing region is concentrated in a narrow ring then there will be no smoothing of the response but the delay can be arbitrarily long, for an arbitrarily large ring radius."
 Evicenth. the spherical shell is the best. bet to smooth out the long-term. variability while producing a short lag. as required. by the data. but we explore. other ecometries too for further analysis.," Evidently, the spherical shell is the best bet to smooth out the long-term variability while producing a short lag, as required by the data, but we explore other geometries too for further analysis."
 We detail cach model in the following subsections., We detail each model in the following subsections.
"value. but slightly higher in the optically Chick regions inside r,= 10 AU.","value, but slightly higher in the optically thick regions inside $r_o =$ 10 AU."
 The deviations [rom steady state are less than156., The deviations from steady state are less than.
" Similarly. the radiative flux relaxes to very close to the analvtical values. except at the center and just inside r,= 10 AU. where a lew cells drop noliceably below the analytical solution."," Similarly, the radiative flux relaxes to very close to the analytical values, except at the center and just inside $r_o =$ 10 AU, where a few cells drop noticeably below the analytical solution."
 To the extent that these few numerical values ciller from the analvlical values. thev err on the side of higher temperatures ancl lower radiative [Iuxes. i.e. (hev err on the side of discouraging cooling of the optically thick regions.," To the extent that these few numerical values differ from the analytical values, they err on the side of higher temperatures and lower radiative fluxes, i.e., they err on the side of discouraging cooling of the optically thick regions."
 Models were also calculated with Eddineton approximation radiative transler and limited diffusion approximation radiative transfer (Boss 2008). testing their ability to maintain the analvtical solution.," Models were also calculated with Eddington approximation radiative transfer and flux-limited diffusion approximation radiative transfer (Boss 2008), testing their ability to maintain the analytical solution."
 The results for these models were identical to those for models with the diffusion approximation. in the latter case because the flux-Iimiter was never triggered bv (he radiative fluxes in the optically thick regions.," The results for these models were identical to those for models with the diffusion approximation, in the latter case because the flux-limiter was never triggered by the radiative fluxes in the optically thick regions."
 In all cases. the code maintains the steady state solution to the same degree that is exhibited by Model R+50 in Figures 5 and 6.," In all cases, the code maintains the steady state solution to the same degree that is exhibited by Model R+50 in Figures 5 and 6."
 Models were also caleulated [or a cloud with the opacity lowered by a [actor of 100. e... &=107 cnr? |.," Models were also calculated for a cloud with the opacity lowered by a factor of 100, i.e., $\kappa = 10^{-2}$ $^2$ $^{-1}$."
 As a result. the central optical depth dropped from ~107 in model R50 to ~107.," As a result, the central optical depth dropped from $\sim 10^4$ in model R+50 to $\sim 10^2$."
 The lower optical depth model was calculated with the Eddington approximation code. ancl ib maintained a temperature profile (hat was only slightly higher (han the analvtical prolile throughout most of disk. again erring on the side of a slightly hotter disk.," The lower optical depth model was calculated with the Eddington approximation code, and it maintained a temperature profile that was only slightly higher than the analytical profile throughout most of disk, again erring on the side of a slightly hotter disk."
 Models were also calculated with and perturbations to the temperature. which behaved similarly to model R+50 except for (heir relaxation to the steady state value occurring faster in time. as would be expected.," Models were also calculated with and perturbations to the temperature, which behaved similarly to model R+50 except for their relaxation to the steady state value occurring faster in time, as would be expected."
 A model was also caleulated with a negative temperature perturbation. Le.. a cloud with the temperature reduced by from + to 5 AU.," A model was also calculated with a negative temperature perturbation, i.e., a cloud with the temperature reduced by from 4 to 5 AU."
 This eloud also relaxed back to the steady state solution within a few Myr. as in model R+50.," This cloud also relaxed back to the steady state solution within a few Myr, as in model R+50."
 Interestingly. in this model the relaxed solution was again slightly hotter than the steady state solution (bv less than 1%)). showing that the code relaxes to this solution independently of heing approached from hotter or cooler temperatures than the steady state solution.," Interestingly, in this model the relaxed solution was again slightly hotter than the steady state solution (by less than ), showing that the code relaxes to this solution independently of being approached from hotter or cooler temperatures than the steady state solution."
 This result implies that the radiative transfer routines. while accurate to a hieh degree. do result in a slight overestimate of cloud temperatures and an uiderestimate of the radiative [luxes. which errs on (he side of discouraging disk instabilitys ability to. produce sell-2ravitating «ΠΗΡΕ.," This result implies that the radiative transfer routines, while accurate to a high degree, do result in a slight overestimate of cloud temperatures and an underestimate of the radiative fluxes, which errs on the side of discouraging disk instability's ability to produce self-gravitating clumps."
 Model T+50 was also calculated with diffusion approximation radiative transfer (Edclington approximation radiative transfer becomes punitivelv slow in multiple dimensions. because of the need to iterate on the solution of the mean intensity equation).," Model T+50 was also calculated with diffusion approximation radiative transfer (Eddington approximation radiative transfer becomes punitively slow in multiple dimensions, because of the need to iterate on the solution of the mean intensity equation)."
" Model T+50 was started with the analvtieal solution for the initial temperature profile for a disk with 9,0=83.1 degrees", Model T+50 was started with the analytical solution for the initial temperature profile for a disk with $\theta_o = 83.1$ degrees
" Model T+50 was started with the analvtieal solution for the initial temperature profile for a disk with 9,0=83.1 degreese", Model T+50 was started with the analytical solution for the initial temperature profile for a disk with $\theta_o = 83.1$ degrees
more stars.,more stars.
 Therefore. while the detectious in IATACS are indeed detections of RGB stars iu the halo of NCC 253. interpretation of the umber counts requires artificial star tests in order to determine the relationship between the numbers aud properties of stars in the IMACS catalog and the true nuubers aud properties of the stellar population.," Therefore, while the detections in IMACS are indeed detections of RGB stars in the halo of NGC 253, interpretation of the number counts requires artificial star tests in order to determine the relationship between the numbers and properties of stars in the IMACS catalog and the true numbers and properties of the stellar population."
 Iu order to test the photometric accuracy and recovery(contamination rates of TRGD stars in NGC 253. we carried out a suite of artificial star tests.," In order to test the photometric accuracy and recovery/contamination rates of TRGB stars in NGC 253, we carried out a suite of artificial star tests."
 Artificial stars with 1126 were placed iuto a lk x lk region iu the outer parts of the Wo aud {- NGC 253 images with empirical PSFs generated roni well-exposed stars., Artificial stars with $m_I < 26$ were placed into a 1k x 1k region in the outer parts of the $V$ and $I$ -band NGC 253 images with empirical PSFs generated from well-exposed stars.
 V-band aud Z-baud appareut uaenitudes were Monte-Carlo sampled frou the CMD of he HST/CGIIOSTS NGC 253 Field 10 pointing., $V$ -band and $I$ -band apparent magnitudes were Monte-Carlo sampled from the CMD of the HST/GHOSTS NGC 253 Field 10 pointing.
" The stars were distributed raudonmly across the region at 8 different deusities (labelled. ""3087 through ""1000k7 depeudiug ou he umber of artificial stars that were added). spanning he observed rauge of RGB star deusitics secu in the IMACS image and in the GIIOSTS data."," The stars were distributed randomly across the region at 8 different densities (labelled “30k” through “4000k” depending on the number of artificial stars that were added), spanning the observed range of RGB star densities seen in the IMACS image and in the GHOSTS data."
 The images were then output aud processed bv the photometric xpeliue in the same way as the cata., The images were then output and processed by the photometric pipeline in the same way as the data.
 Photometry was performed ou the artificial star frames exactly as for the original frames., Photometry was performed on the artificial star frames exactly as for the original frames.
 CMDs generated from he input magnitudes aud from the output photometry ou the artificial star frames forthree sample artificial star tests (the 30k. τοῦ aud LOOOk tests) are shown in Figure L.," CMDs generated from the input magnitudes and from the output photometry on the artificial star frames forthree sample artificial star tests (the $30$ k, $700$ k and $4000$ k tests) are shown in Figure \ref{fig:artcmd}."
 As the stellar density increases. the ROB stars become blended anc are detected at maguitucdes sieuificautlv brighter than those of mdividual stars.," As the stellar density increases, the RGB stars become blended and are detected at magnitudes significantly brighter than those of individual stars."
 This results in a CAID exactly like the oue observed if the niajoritv of our sources are found in regions of high deusitv., This results in a CMD exactly like the one observed if the majority of our sources are found in regions of high density.
 Tn order to determine the παπανο density of stars at the tip of the red giant branch (TROB). we have compared the number of input RGB stars above a threshold of I«21.25 (correspouding to an absolute maguitude of Mj<3.16 at the distance of NGC 253) to the mmuber density of output detections within a box extending frou I—23 to 21. and from V7=1.0 to 2.5.," In order to determine the number density of stars at the tip of the red giant branch (TRGB), we have compared the number of input RGB stars above a threshold of $I < 24.25$ (corresponding to an absolute magnitude of $M_I < -3.46$ at the distance of NGC 253) to the number density of output detections within a box extending from $I=23$ to $24$, and from $V-I=1.0$ to $2.5$."
 This box lies entirely above the input threshold iu order to better capture the bulk of the halo stars at the typical deusities of our data (see Figure [))., This box lies entirely above the input threshold in order to better capture the bulk of the halo stars at the typical densities of our data (see Figure \ref{fig:artcmd}) ).
 We lave tested several ifereut criteria for both the TRGB threshold aud the magnitude boundaries of the selection box. raugiug bv £0.5 mag.," We have tested several different criteria for both the TRGB threshold and the magnitude boundaries of the selection box, ranging by $\pm 0.5$ mag."
 We fiud that this introduces ~20% changes in the derived iunuber deusities once they are normalized bv the expected muuber of stars above each TRGB threshold for an old ietal-poor population. as discussed in Section 5..," We find that this introduces $\sim 20\%$ changes in the derived number densities once they are normalized by the expected number of stars above each TRGB threshold for an old metal-poor population, as discussed in Section \ref{sec:haloluminosity}."
 The relationship between the iuput nunuber density ove the TRGD threshold aud the output nuuuber density withiu the sclection box is shown as the right-hand (high-deusitv) dashed line in the upper panel of Figure 5.. while their ratio (ic. the recovered fraction) is shown in the bottom panel.," The relationship between the input number density above the TRGB threshold and the output number density within the selection box is shown as the right-hand (high-density) dashed line in the upper panel of Figure \ref{fig:artrecov}, while their ratio (i.e. the recovered fraction) is shown in the bottom panel."
 A problem is imuueciately apparent: because there are 183 detections within the RGB selection box iu the “blank field” to which the artificial stars were added (a combination of background ealaxies aud real ROB stars in the stellar halo). more stars are detected than were added at low clensitics.," A problem is immediately apparent: because there are $183$ detections within the RGB selection box in the “blank field” to which the artificial stars were added (a combination of background galaxies and real RGB stars in the stellar halo), more stars are detected than were added at low densities."
 This is a property of the baseline nuage. not of our ability o detect stars. and so is artificial.," This is a property of the baseline image, not of our ability to detect stars, and so is artificial."
 One could simply subtract 183 from the wuuber of output stars (shown as he left-handed dashed line iu Figure 5)). but this is also incorrect: at hieh deusities. we approach the confusion inited reginae. aud the original detections have mostly con covered up by artificial stars ancl are therefore no onger detected.," One could simply subtract $183$ from the number of output stars (shown as the left-handed dashed line in Figure \ref{fig:artrecov}) ), but this is also incorrect: at high densities, we approach the confusion limited regime, and the original detections have mostly been covered up by artificial stars and are therefore no longer detected."
 A better approach is to compare the ocations of the detections iu the artificial star frames to he locations of the original detections in the blank field., A better approach is to compare the locations of the detections in the artificial star frames to the locations of the original detections in the blank field.
 We pertormed a elobally-optimized one-to-one match of he closest original detection to cach detection in the artificial star frame. aud omitted auv whose best match was within 075.," We performed a globally-optimized one-to-one match of the closest original detection to each detection in the artificial star frame, and omitted any whose best match was within $0\farcs 5$."
 This results in the black solid line in Figure 5.., This results in the black solid line in Figure \ref{fig:artrecov}.
 The fraction of blank Ποια sourcees that are onütted. fusus. IX a useful measure of the deeree to which individual sources are likely to be covered up.," The fraction of blank field sourcees that are omitted, $f_{\mathrm{bg,omit}}$, is a useful measure of the degree to which individual sources are likely to be covered up."
 At the confusion limit. the uuuber of output detections saturates. and even begins to decrease as the majority of the blends become brighter than the selection box.," At the confusion limit, the number of output detections saturates, and even begins to decrease as the majority of the blends become brighter than the selection box."
" We therefore set a threshold of logs=5.60. at which we can only set a lower limit ou the true deusitv (log9g,zx 6.39). shown as the right-haucl vertical eray lue in Figure 5.."," We therefore set a threshold of $\log n_{\mathrm{out}}=5.60$, at which we can only set a lower limit on the true density $\log n_{\mathrm{in}} \ge 6.39$ ), shown as the right-hand vertical gray line in Figure \ref{fig:artrecov}. ."
 Fortunately. very little of the dat: reach this lait.," Fortunately, very little of the data reach this limit."
 At the low density cud. we note that the recovered fraction approaches unity. aud so we male," At the low density end, we note that the recovered fraction approaches unity, and so we make"
along the local largescale structure. and thus similarly to cach other.,"along the local large–scale structure, and thus similarly to each other."
 Ou larger scales. such elougation iu cosmic structures appears to be preseut.," On larger scales, such elongation in cosmic structures appears to be present."
 For instance. the effects of large scale structure ou the shapes of galaxy clusters have been the subject of much. study.," For instance, the effects of large scale structure on the shapes of galaxy clusters have been the subject of much study."
 Cosmological Nbody simulations lave sugeested that clusters tend to be oricuted towards ucighboring clusters or in clirections defined by adjoining filaments aud the mereine subclusters which drain aloug them. (, Cosmological N--body simulations have suggested that clusters tend to be oriented towards neighboring clusters or in directions defined by adjoining filaments and the merging subclusters which drain along them. (
Dekel. West Aarscth 1981: West. Dekel Οσο 1989: West. Vilhuusen Dekel 1991: van Waarlem van de Weveacrt 1993: Splinter 1997: de Theije. van Iampen Slijkhuis 1995).,"Dekel, West Aarseth 1984; West, Dekel Oemler 1989; West, Villumsen Dekel 1991; van Haarlem van de Weygaert 1993; Splinter 1997; de Theije, van Kampen Slijkhuis 1998)."
" Observations have typically indicated the presence of such nueuts. either towards nearby clusters (Binegech 1982: Flin 1987: West L989a.b: Rhee. van Haarlem Kateert 1992: Pliouis 199L) or towards nearby large.scale structure in the ealaxy distribution (Arevres 1986: μάνας, Groth Peebles 1988): although not all studies support the presence of such aligunmieuts (Struble Peebles 1985: Ulmer. MeMillan. Isowalski 1989: Fone. Stevenson Shanks 1990)."," Observations have typically indicated the presence of such alignments, either towards nearby clusters (Binggeli 1982; Flin 1987; West 1989a,b; Rhee, van Haarlem Katgert 1992; Plionis 1994) or towards nearby large–scale structure in the galaxy distribution (Argyres 1986; Lambas, Groth Peebles 1988); although not all studies support the presence of such alignments (Struble Peebles 1985; Ulmer, McMillan Kowalski 1989; Fong, Stevenson Shanks 1990)."
 The existence ofcorrelations in the aliguineut of largescale structures appears quite possible: perhaps simular iutriusic correlations in aliguimcut exist ou galactic scales., The existence of correlations in the alignment of large–scale structures appears quite possible; perhaps similar intrinsic correlations in alignment exist on galactic scales.
 Theoretical expectations for the degree of correlation of intrinsic cllipticities ean in principle be derived., Theoretical expectations for the degree of correlation of intrinsic ellipticities can in principle be derived.
 The local eravitational shear can be expected to either align the Intrinsic aneular momentum of nearby galaxies (Lee Peu 2000). or to sinübulv deform neiehboriusg. nourotating ealaxies through tidal distortion (Ciotti Dutta 1991).," The local gravitational shear can be expected to either align the intrinsic angular momentum of nearby galaxies (Lee Pen 2000), or to similarly deform neighboring, non–rotating galaxies through tidal distortion (Ciotti Dutta 1994)."
 Thus. the statistics of the local tidal field can be related to the statistics of galaxy augular momenta (Catclan Theuns 1996a.b: Catelan Theuns 1997: Suecriman. Sunuucrs Iunionunkowski 1999): and therefore to the intrinsic correlations in galaxy cllipticitics (Coutts 1996: Lee Pen 2000: Catelan 2000. in preparation: Mackey aud White 2000. in preparation).," Thus, the statistics of the local tidal field can be related to the statistics of galaxy angular momenta (Catelan Theuns 1996a,b; Catelan Theuns 1997; Sugerman, Summers Kamionkowski 1999); and therefore to the intrinsic correlations in galaxy ellipticities (Coutts 1996; Lee Pen 2000; Catelan 2000, in preparation; Mackey and White 2000, in preparation)."
 There have been nuuucrous attempts to detect intrinsic correlations in galaxy aliguiucuts using low redslüft samples: he picture painted by this work is unclear. as we cau see frou the following sample.," There have been numerous attempts to detect intrinsic correlations in galaxy alignments using low redshift samples; the picture painted by this work is unclear, as we can see from the following sample."
 Flin (1988) considered a suuple of 11δ ealaxics in the Perseus supercluster aud ound that the spin axes of these galaxies were aligned with he supercluster plane., Flin (1988) considered a sample of 118 galaxies in the Perseus supercluster and found that the spin axes of these galaxies were aligned with the supercluster plane.
 Muriel Lanmbas (1992) reported a correlation of aligumuents seen with spirals taken frou he ESO catalog and analyzed in three dimensions: whew ilv projected data was considered. the correlation was 10 longer preseut.," Muriel Lambas (1992) reported a correlation of alignments seen with spirals taken from the ESO catalog and analyzed in three dimensions; when only projected data was considered, the correlation was no longer present."
 Garrido (1993) analyzed a sample covering a large area of sky in the northern hemisphere aud claimed to find no evidence for correlations in aliguineut except witlin the Coma supercluster., Garrido (1993) analyzed a sample covering a large area of sky in the northern hemisphere and claimed to find no evidence for correlations in alignment except within the Coma supercluster.
 Tan. Gould Sackett (1995) examined the spins of 60 galaxies in the Ursa Major filament aud found no evidence for any aliguinent of spins.," Han, Gould Sackett (1995) examined the spins of 60 galaxies in the Ursa Major filament and found no evidence for any alignment of spins."
 Cabanela Aldering (1998) considered galaxy shapes xtracted frou α survey of Perseus.Pisces couducted usiug ji automated plate scanner: statistically sienificaut aud color dependent correlatious of galaxy cllipticities were found., Cabanela Aldering (1998) considered galaxy shapes extracted from a survey of Perseus–Pisces conducted using an automated plate scanner; statistically significant and color dependent correlations of galaxy ellipticities were found.
 Ou the other hand. Cabanela Dickey (1999) used III observations to determine the spins of 51 galaxies in the PerseusPisces supercluster: and found no evidence for preferential aligmmeuts of spin vectors.," On the other hand, Cabanela Dickey (1999) used HI observations to determine the spins of 54 galaxies in the Perseus–Pisces supercluster; and found no evidence for preferential alignments of spin vectors."
 At this time evideuce favors au orieutatiou aliguinent between eD galaxies and the major axis of their pareut cluster: but the preseuce or absence of auy other galaxy shape correlations remiss undetermined., At this time evidence favors an orientation alignment between cD galaxies and the major axis of their parent cluster; but the presence or absence of any other galaxy shape correlations remains undetermined.
 Iu this paper. we use Nbody simulations to make theoretical xedietious for the correlation of iutrinsie galaxy ellipticities.," In this paper, we use Nbody simulations to make theoretical predictions for the correlation of intrinsic galaxy ellipticities."
 For simplicity. we work with directly with the projected ellipticities of the simulated dark matter halos. without naking assuming auv model for the wav galaxies form within them.," For simplicity, we work with directly with the projected ellipticities of the simulated dark matter halos, without making assuming any model for the way galaxies form within them."
 The significance of our results will therefore )o entirely dependent on whether galaxy cllipticitics behave sienificantly differentle from thei halos. a problem we cave to future easdvuaiical sinmlations.," The significance of our results will therefore be entirely dependent on whether galaxy ellipticities behave significantly differently from their halos, a problem we leave to future gasdynamical simulations."
 The lavout of the paper is as follows., The layout of the paper is as follows.
 Iu 822. we describe IGN body dataset used. our halo catalog. aud our 1ieasureimoenut f projected. ellipticities from the halos.," In 2, we describe the N–body dataset used, our halo catalog, and our measurement of projected ellipticities from the halos."
 We measure the iree dinieusioual correlation functious of projected ellipticities oe1 833. and in 811 we decribe the coustruction of simulated survevs from our halo catalogs. with a geometry desieued o nundc weak lensing observations.," We measure the three dimensional correlation functions of projected ellipticities in 3, and in 4 we decribe the construction of simulated surveys from our halo catalogs, with a geometry designed to mimic weak lensing observations."
 Tn 855. we project ie three dinensioual correlation functions iuto aneular y.atistics. includiug the shear variance.," In 5, we project the three dimensional correlation functions into angular statistics, including the shear variance."
 We also compute jese aneular measures cirectlv from our simulated surveys (as a consistency clieck)., We also compute these angular measures directly from our simulated surveys (as a consistency check).
 We compare our results to current observational data iu 866. before discussing aud σπασαι our results iu 877.," We compare our results to current observational data in 6, before discussing and summarizing our results in 7."
 Our requirements are that the simulatiou volume be aree enough to capture the large scale tidal field that may cause correlatious to arise between halo shapes. while at he same time having enough mass resolution to follow the ormation of galaxy sized halos with a reasonable ΠΡΟΣ of particles.," Our requirements are that the simulation volume be large enough to capture the large scale tidal field that may cause correlations to arise between halo shapes, while at the same time having enough mass resolution to follow the formation of galaxy sized halos with a reasonable number of particles."
 We use outputs frou Nbody simulations run wo the Virgo Consortimm (see ce. the author list of Jenkins 1998 for Virgo members) aud which they ive generously made public.," We use outputs from Nbody simulations run by the Virgo Consortium (see e.g., the author list of Jenkins 1998 for Virgo members) and which they have generously made public."
 The simulations are part of a set of different models. although we only use one here. he currently favoured cosmological constaut-«donünatec cold dark matter (ACDAL) model.," The simulations are part of a set of different models, although we only use one here, the currently favoured cosmological constant-dominated cold dark matter $\Lambda$ CDM) model."
 The parameters of this nodel are as follows. ο=0.3 (the matter denstv at = (n. 94=0.7 (the contribution of A to the critica density at z=0). P=0.21. which is the shape parameter of the initial power spectrmu.," The parameters of this model are as follows, $\Omega_{m}=0.3$ (the matter density at $z=0$ ), $\Omega_{\Lambda}=0.7$ (the contribution of $\Lambda$ to the critical density at z=0), $\Gamma=0.21$, which is the shape parameter of the initial power spectrum."
 The normalization was set so that oy= 0.9. (the rs matter fluctuations in 8fi1Mpe spheres extrapolated to += 0).," The normalization was set so that $\sigma_{8}=0.9$ , (the rms matter fluctuations in $8 \hmpc$ spheres extrapolated to $z=0$ )."
 We use two different simulations. one un with a box size ofLIL35.tMpc. and another iu a box of size 2105. 1Mpe. both with 256? particles.," We use two different simulations, one run with a box size of $141.3 \hmpc$, and another in a box of size $240 \hmpc$ , both with $256^3$ particles."
 The mass per particle in the former case is Lob&LOMbLTALL. aud iu the second it is 5 Έπος larger.," The mass per particle in the former case is $1.4 \times 10^{10} \msun$, and in the second it is 5 times larger."
 The Nbody code used was au adaptive particle-particle particleauesh code (Couchiman. Thomas and Pearce). and the eravitational softening length was 30hUspe (for the smaller box).," The Nbody code used was an adaptive particle-particle particle-mesh code (Couchman, Thomas and Pearce), and the gravitational softening length was $30 \hkpc$ (for the smaller box)."
 For more details. the reacler is referred to the Virgo papers. including Ἱναιια (1999) for the sumaller box simulation.," For more details, the reader is referred to the Virgo papers, including Kauffmann (1999) for the smaller box simulation."
 Unless stated otherwise. all our analysis will be carried out using the lugher resolution simulation. with the other being used as a check of the effects of lower resolution.," Unless stated otherwise, all our analysis will be carried out using the higher resolution simulation, with the other being used as a check of the effects of lower resolution."
 We use only the 2=] output in each case. as we will be interested iu comparing to lensing obscrvatious where the peaks of the ealaxy distributionis expected to lieclose to this redshift.," We use only the $z=1$ output in each case, as we will be interested in comparing to lensing observations where the peak of the galaxy distributionis expected to lieclose to this redshift."
Corouagraphs that could in theory be cousidered for this dual-mask approach include uearly every Lyot-type coronagraph with a racially-svuuuetric local-plane mask. iucludiug Lyot coronagraphis. APLCs. BLCs with ractially-svunmetric masks (?).. LOPMs. and VCs.,"Coronagraphs that could in theory be considered for this dual-mask approach include nearly every Lyot-type coronagraph with a radially-symmetric focal-plane mask, including Lyot coronagraphs, APLCs, BLCs with radially-symmetric masks \mycitep{Kuc03}, 4QPMs, and VCs."
 Uufortunately. this simple analysis ignores one large difficulty: not only does mask A block the core of star A. but mask B will suppress a portion of the sidelobes of star A. which then will scatter light across the image plane. limiting the coutrast that can be achieved.," Unfortunately, this simple analysis ignores one large difficulty: not only does mask A block the core of star A, but mask B will suppress a portion of the sidelobes of star A, which then will scatter light across the image plane, limiting the contrast that can be achieved."
 Thus. ouly a coronagraph with the lowest sidelobes at the first image plane is likely to prove promising lor a binary system.," Thus, only a coronagraph with the lowest sidelobes at the first image plane is likely to prove promising for a binary system."
 We investigate the APLC because it includes an apocization at the entrance pupil that reshapes the PSF., We investigate the APLC because it includes an apodization at the entrance pupil that reshapes the PSF.
 As a result. the APLC has optimally low sidelobes (?) aud will minimize the interaction between the two masks.," As a result, the APLC has optimally low sidelobes \mycitep{Sou03} and will minimize the interaction between the two masks."
 Most existing telescopes have obstructed apertures. and the effects of secondaries aud their support structures on the coronagraplis must be considered iu the mocdeliug of the optical performance.," Most existing telescopes have obstructed apertures, and the effects of secondaries and their support structures on the coronagraphs must be considered in the modeling of the optical performance."
 In many cases. the effects of these optical obstructions cau be mitigated by carefully. desiguiug the coronagrapl," In many cases, the effects of these optical obstructions can be mitigated by carefully designing the coronagraph."
 For example. Lyot stops cau be designed to block the diffractiou [fromaperture obstructions (?).. spider removal plates can remove the spiders [rom the image (?).. and for APLCs. the apodization can be optimized for arbitrary apertures (?)..," For example, Lyot stops can be designed to block the diffraction fromaperture obstructions \mycitep{Siv05}, spider removal plates can remove the spiders from the image \mycitep{Loz09}, and for APLCs, the apodization can be optimized for arbitrary apertures \mycitep{Sou09}."
 However. the reduced performance should be taken into account when cesignius tliese coronagraplis.," However, the reduced performance should be taken into account when designing these coronagraphs."
 ? makes binary star observations with the unobstructed subaperture (?) on the Palomar Observatorys 200-inel Hale Telescope. neatly avoiding these coronagraphic complications. but ouly utilizing 95€ of the collecting area of the telescope.," \mycitet{Cre10} makes binary star observations with the unobstructed subaperture \mycitep{Ser07} on the Palomar Observatory's 200-inch Hale Telescope, neatly avoiding these coronagraphic complications, but only utilizing $\sim9\%$ of the collecting area of the telescope."
 A further difficulty. arises [rom the field rotation. which causes the binary system to rotate with respect to the telescope pupil [or telescopes in altitude-aziuuth mouuts.," A further difficulty arises from the field rotation, which causes the binary system to rotate with respect to the telescope pupil for telescopes in altitude-azimuth mounts."
 In high contrast observations. this characteristic can be employed for Aneular Differential Imagine (ADI) (?," In high contrast observations, this characteristic can be employed for Angular Differential Imaging (ADI) \mycitep{Mar06}."
 The major cousequence of this is that the coronagraphi must be able to maintain its performance regardless of the orientation of the aperture with respect to the image plaue mask. includingD> any obstructious.," The major consequence of this is that the coronagraph must be able to maintain its performance regardless of the orientation of the aperture with respect to the image plane mask, including any obstructions."
 For an unobstructed circular aperture. the circular symmetry of the pupil reuders this trivial.," For an unobstructed circular aperture, the circular symmetry of the pupil renders this trivial."
 The performance of the system is also depeudeut ou the aberrations introduced by the atiuosphliere aud the ability of the adaptive optics CAO) system to correct these wavefront errors., The performance of the system is also dependent on the aberrations introduced by the atmosphere and the ability of the adaptive optics (AO) system to correct these wavefront errors.
 To examine the performance of the coronagraphic optious discussed above. we simulate their performance on a realistic system.," To examine the performance of the coronagraphic options discussed above, we simulate their performance on a realistic system."
 Here. we use parameters of the Subaru Telescope aud HICLAO (?) [or simulation purposes: the pupil is a 5.211 diameter circular aperture with au IR secoucdary 1.265m diameter circular obscuration coucentric with the primary mirror (?)..," Here, we use parameters of the Subaru Telescope and HiCIAO \mycitep{Hod08} for simulation purposes: the pupil is a $8.2$ m diameter circular aperture with an IR secondary $1.265$ m diameter circular obscuration concentric with the primary mirror \mycitep{Usu03}."
 The secondary is held in place with [our spiders. as shown in 2.," The secondary is held in place with four spiders, as shown in ."
. From a numerical perspective. we make extensive," From a numerical perspective, we make extensive"
galaxy’s local environment.,galaxy's local environment.
 We therefore limit our sample to the subset of SLACS lenses with published inner density slopes., We therefore limit our sample to the subset of SLACS lenses with published inner density slopes.
 This saniple consists of 15 earlv-tv galaxies with lens redshifts between z=0.063 and z=0.332pe and stellar velocity. dispersions from ao=178 to 0=330 (IKoopmaunsetal.2006)., This sample consists of 15 early-type galaxies with lens redshifts between $z = 0.063$ and $z = 0.332$ and stellar velocity dispersions from $\sigma = 178$ to $\sigma = 330$ \citep{koopmansSLACS}.
" The density slopes range [rom a= Ls2toa22.34 where pxr. ""and the typical error on a for a given system is approximately 5 per cent (IXoopmiansetal. 2006)."," The density slopes range from $\alpha = 1.82$ to $\alpha = 2.34$ where $\rho \propto r^{-\alpha}$, and the typical error on $\alpha$ for a given system is approximately 5 per cent \citep{koopmansSLACS}."
. Four of the systems. have shallower than isothermal density slopes (Le. a« 2). six are approximately consistent with isothermal. and five lenses have steeper than isothermal profiles.," Four of the systems have shallower than isothermal density slopes (i.e., $\alpha < 2$ ), six are approximately consistent with isothermal, and five lenses have steeper than isothermal profiles."
 Phe lens characteristics are summarized in Table 1.., The lens characteristics are summarized in Table \ref{table_lens_properties}.
 Photometric and spectroscopic data from the SDSS Data Release 6 (Acelman-MeCarthyctal.2007). are. used to quantify the environments. of the lensing galaxics., Photometric and spectroscopic data from the SDSS Data Release 6 \citep{sdss} are used to quantify the environments of the lensing galaxies.
 Catalogs are mace of all SDSS primary galaxies with proc22.1] mag that lie within L2 ‘Alpe of the lensing galaxy., Catalogs are made of all SDSS primary galaxies with $r < 22.1$ mag that lie within 1.2 Mpc of the lensing galaxy.
 These catalogues include the [ive banc SDSS photometry. photometric redshifts from the table of the SDSS catalogue. ancl spectroscopic redshifts for those ealaxies targeted by SDSS.," These catalogues include the five band SDSS photometry, photometric redshifts from the table of the SDSS catalogue, and spectroscopic redshifts for those galaxies targeted by SDSS."
 Due to the range in redshifts spanned by the lenses. the spectroscopic completeness varies significantly from svstem to system.," Due to the range in redshifts spanned by the lenses, the spectroscopic completeness varies significantly from system to system."
 Only the svstenmis with z01 contain sullicient spectroscopic redshifts to adequately. characterize the elobal environments of the lenses., Only the systems with $z \la 0.1$ contain sufficient spectroscopic redshifts to adequately characterize the global environments of the lenses.
 We use the colours of the field. galaxies and their ollsets from the lens to determine the likelihood that a neighbour galaxy is physically associated with the lens., We use the colours of the field galaxies and their offsets from the lens to determine the likelihood that a neighbour galaxy is physically associated with the lens.
 Only galaxies that are less than one magnitude brighter or 2.5 magnitucles fainter than the lens are used in the analysis., Only galaxies that are less than one magnitude brighter or 2.5 magnitudes fainter than the lens are used in the analysis.
 Piclucial colours of galaxies at the lens redshift are determined by. querving all spectroscopic galaxies in the SDSS database that have redshifts within 1200 of the redshift of the lensing galaxy (this is equivalent to Aszz0.004 at the redshifts of the lenses)., Fiducial colours of galaxies at the lens redshift are determined by querying all spectroscopic galaxies in the SDSS database that have redshifts within 1200 of the redshift of the lensing galaxy (this is equivalent to $\Delta z \approx 0.004$ at the redshifts of the lenses).
 These empirical colour distributions are used to determine Gaussian. weights for cach of the field. galaxies. which are also weighted by their distance [rom the lens system.," These empirical colour distributions are used to determine Gaussian weights for each of the field galaxies, which are also weighted by their distance from the lens system."
 The colour distributions for a 2=0.25 galaxy are shown in Figure 1: note that there are. long tails at the blue end. of cach clistribution., The colour distributions for a $z = 0.25$ galaxy are shown in Figure \ref{figure_color_distribution}; note that there are long tails at the blue end of each distribution.
 These tails are due to late-tvpe galaxies ancl are elipped from our analysis when determining the location and width of the distributions., These tails are due to late-type galaxies and are clipped from our analysis when determining the location and width of the distributions.
 The clipping biases us against finding late-type companion ealaxies: however. this is not a strong bias because carly-tvpe galaxies are known to cluster much more strongly than Iate-type galaxies (Meneuxetal.2006).," The clipping biases us against finding late-type companion galaxies; however, this is not a strong bias because early-type galaxies are known to cluster much more strongly than late-type galaxies \citep{menaux}."
. The final weighting scheme used is where ρω is the weight for cach galaxy ancl f(d) is eiven by where d is the distance of the galaxy from the lens and dy is set to 20 tkpe for an elfective distance. of approximately 100 tkpe (Le.severaltruncationradii.Limousinetal.," The final weighting scheme used is where $w_{gal}$ is the weight for each galaxy and $f(d)$ is given by where $d$ is the distance of the galaxy from the lens and $d_0$ is set to 20 kpc for an effective distance of approximately 100 kpc \citep[i.e. several truncation radii,][]{limousin}."
" 2007).. The product is taken over the set of colours e=[g|rrLioios ds the dillerence between the mean of the empirical colour distribution and the measured colour of the galaxy. and a, is the quadrature sum of the width of the colour distribution and the errors on the SDSS photometry."," The product is taken over the set of colours $c =\{g-r, r-i, i-z\}$, $\Delta_{c}$ is the difference between the mean of the empirical colour distribution and the measured colour of the galaxy, and $\sigma_{c}$ is the quadrature sum of the width of the colour distribution and the errors on the SDSS photometry."
 The a filter is not used in our analysis due to its poor sensitivity ancl all of the galaxies in our sample have the three color photometry {gqrer£z] available from the SDSS database.," The $u$ filter is not used in our analysis due to its poor sensitivity and all of the galaxies in our sample have the three color photometry $\{g-r, r-i, i-z\}$ available from the SDSS database."
" We sum the weights of all of the galaxies in the field to determine an effective number of galaxies physically associated with the lensing galaxy. iN,"," We sum the weights of all of the galaxies in the field to determine an effective number of galaxies physically associated with the lensing galaxy, $N_w$."
 The photometric weighting is not equivalent to determining a photometric redshift for cach galaxy: photometric redshifts characterize the most likely redshift of a source while our weights are a proxy for the likelihood of a source being at a given redshift (though the full probability clistribution from a photometric redshift analysis could approximately provide the same information)., The photometric weighting is not equivalent to determining a photometric redshift for each galaxy; photometric redshifts characterize the most likely redshift of a source while our weights are a proxy for the likelihood of a source being at a given redshift (though the full probability distribution from a photometric redshift analysis could approximately provide the same information).
 Our weighting algorithm was tested on. low-reclshilt (0.025<2« 0.075) SDSS galaxies that have Iuminosities and stellar. velocity. dispersions that are. comparable. to the SLACS lenses., Our weighting algorithm was tested on low-redshift $0.025 < z < 0.075$ ) SDSS galaxies that have luminosities and stellar velocity dispersions that are comparable to the SLACS lenses.
 Phese low-redshift galaxies have nearly complete spectroscopy for. all neighbouring field. galaxies with rroen€ggd2.5., These low-redshift galaxies have nearly complete spectroscopy for all neighbouring field galaxies with $r_{field} < r_{gal} + 2.5$.
" The weighting algorithm was emploved as above. and a comparison was made between ;V,, and the number of spectroscopically confirmed. neighbours."," The weighting algorithm was employed as above, and a comparison was made between $N_w$ and the number of spectroscopically confirmed neighbours."
" ANS correctly estimated. thefofaé number of companions within 100 !kpe of the target galaxy for 64 per cont of the systems. where IN, was rounded to the nearest integer to make the comparison."," $N_w$ correctly estimated the number of companions within 100 kpc of the target galaxy for 64 per cent of the systems, where $N_w$ was rounded to the nearest integer to make the comparison."
" For the incorrectly matched systems. N, does not show a bias for over or underestimating the number of companions: LS per cent were incorrectly matched in either case."," For the incorrectly matched systems, $N_w$ does not show a bias for over or underestimating the number of companions; 18 per cent were incorrectly matched in either case."
 However. Ny. is more robust as a binary," However, $N_w$ is more robust as a binary"
 and coutribution from Davis in these proceedings)., and contribution from Davis in these proceedings).
 A distinenishine aspect of DEEP is that the survey ais to gather not oulv very fait redshifts. but also internal kinematic data in the form of rotation curves or line widths. as well as line strengths seusitive to star formation rates. eas conditions. age.o aud moetallicity.," A distinguishing aspect of DEEP is that the survey aims to gather not only very faint redshifts, but also internal kinematic data in the form of rotation curves or line widths, as well as line strengths sensitive to star formation rates, gas conditions, age, and metallicity."
 While waiting for the completion of DEINOS so that the major DEEP survey of ! ealaxies cau beein (see Davis coutributiou). we have becu undertaking a number of smaller. pilot-stvle projects with the existing Low Resolution huagiug Spectrograph (LRIS: [10])) to determine what is feasible with Keck and thus to help refine the scope of the main DEEP survey.," While waiting for the completion of DEIMOS so that the major DEEP survey of $^+$ galaxies can begin (see Davis contribution), we have been undertaking a number of smaller, pilot-style projects with the existing Low Resolution Imaging Spectrograph (LRIS: \cite{Oke95}) ) to determine what is feasible with Keck and thus to help refine the scope of the main DEEP survey."
 To maximize the scientific returns for our relatively stuall samples (currently over 500 galaxies). we observed fields where TST WFPC2 mages already. exist. including the TIDF aud flanking fields [7].. |12].. [2].. |ü]:: the Croth Survey Strip |5].. [0].. and Selected Area 68.," To maximize the scientific returns for our relatively small samples (currently over 500 galaxies), we observed fields where HST WFPC2 images already exist, including the HDF and flanking fields \cite{Low97}, \cite{Phi97}, \cite{Guz97}, \cite{Vogt97}; ; the Groth Survey Strip \cite{Koo96}, \cite{Vogt96}, and Selected Area 68."
 Such HST images provide morphology data aud also the structure. size. aud inclinatiou data needed to couvert kinematic observations from Weck iuto direct niieasures of mass.," Such HST images provide morphology data and also the structure, size, and inclination data needed to convert kinematic observations from Keck into direct measures of mass."
 Our data is still largely being reduced. but we have already achieved a redshift completeness of for a 200 ealaxy sample reaching a limit ofJ—23.5.," Our data is still largely being reduced, but we have already achieved a redshift completeness of for a 200 galaxy sample reaching a limit of $I \sim 23.5$."
 Overall. our fiudiugs have reassured us that our major DEIMOS survey is not only feasible. but that kineiiaties will indeed be a powerful additional dimension of study.," Overall, our findings have reassured us that our major DEIMOS survey is not only feasible, but that kinematics will indeed be a powerful additional dimension of study."
 Our work with line strengths is not vet mature enoueh to be preseuted. so the following will focus ou highlights of our kinematic survevs.," Our work with line strengths is not yet mature enough to be presented, so the following will focus on highlights of our kinematic surveys."
 As seen in Fig., As seen in Fig.
 l. we have clearly demonstrated that emission-line rotation curves of likely spirals can be observed to redshifts near 2~1 for galaxies as faint as [~22 with one to two hour exposures |0]..," 1, we have clearly demonstrated that emission-line rotation curves of likely spirals can be observed to redshifts near $z
\sim 1$ for galaxies as faint as $I \sim 22$ with one to two hour exposures \cite{Vogt96}."
 Based on 16 galaxies so far. we find little evidence for auy major chauge (<0.6 mae) in the zero-poiut of the optical Tully-Fisher relation |0]..," Based on 16 galaxies so far, we find little evidence for any major change $<0.6$ mag) in the zero-point of the optical Tully-Fisher relation \cite{Vogt97}."
 These results are iu stark coutrast to claims for more exteusive evolution of 1.5 mag to 2.0 mag [13].. |L1]. for very blue galaxies.," These results are in stark contrast to claims for more extensive evolution of 1.5 mag to 2.0 mag \cite{Rix97}, \cite{Sim98} for very blue galaxies."
 Larger sauples will be needed to uuderstaud the causes (6.9.. DIuniuositv or color) of these differences.," Larger samples will be needed to understand the causes (e.g., luminosity or color) of these differences."
" ""Though rotation curves are preferable. the vast majority of very faint galaxies are too small to vield more than line widths as kincmatic data."," Though rotation curves are preferable, the vast majority of very faint galaxies are too small to yield more than line widths as kinematic data."
 Except for very bright ealaxies that night vield absorption line widths. emission lines are used.," Except for very bright galaxies that might yield absorption line widths, emission lines are used."
" ""Though winds. dust obscuration.and poorrepreseutation of the eravitational"," Though winds, dust obscuration,and poorrepresentation of the gravitational"
"The Planck observation is hindered by the clustering problem caused by its large PSF (53, rendering its flux estimates completely useless unless a correction is applied.","The Planck observation is hindered by the clustering problem caused by its large PSF (5'), rendering its flux estimates completely useless unless a correction is applied."
" The problem is clearly illustrated in Fig. 14,,"," The problem is clearly illustrated in Fig. \ref{fig:Swire-Planck-Histo},"
 where we show the histograms of the ratio of the flux estimates to the input fluxes for a Planck observation of the SWIRE fields for two selected redshift bins., where we show the histograms of the ratio of the flux estimates to the input fluxes for a Planck observation of the SWIRE fields for two selected redshift bins.
" We developed a simple method to correct this When stacking sources in a given redshift bin with Planck, we measure the added contribution of the sources and the clustering."," We developed a simple method to correct this When stacking sources in a given redshift bin with Planck, we measure the added contribution of the sources and the clustering."
" To correct the stacked fluxes with Planck for the effects of clustering, we use source fluxes at 850 um obtained by stacking SCUBA-2 data."," To correct the stacked fluxes with Planck for the effects of clustering, we use source fluxes at 850 $\mu$ m obtained by stacking SCUBA-2 data."
 If we stack sources detected by Planck for which we have an, If we stack sources detected by Planck for which we have an
is effective within its envelope. is still rather uncertain.,"is effective within its envelope, is still rather uncertain."
 In thisLetter we have shown that late energy injection. from several hundreds of seconds to several hours. is à common feature of a collapsar model in which the pre-supernova star has unmixed shells of burning elements.," In this we have shown that late energy injection, from several hundreds of seconds to several hours, is a common feature of a collapsar model in which the pre-supernova star has unmixed shells of burning elements."
 The lack of flaring activity in more than half of the well-monitored GRBs hence argues against such a model., The lack of flaring activity in more than half of the well-monitored GRBs hence argues against such a model.
 The arguments presented here give support to a fully mixed pre-supernova. as expected for a rapidly-rotating. low-metallicity massive star.," The arguments presented here give support to a fully mixed pre-supernova, as expected for a rapidly-rotating, low-metallicity massive star."
 We thank Abe Falcone. Bing Zhang and Weiqun Zhang for discussions on various aspects of this work.," We thank Abe Falcone, Bing Zhang and Weiqun Zhang for discussions on various aspects of this work."
We thank Ramesh Narayan. Eliot. Quataert. anc Chris Revnolds for useful discussions.,"We thank Ramesh Narayan, Eliot Quataert and Chris Reynolds for useful discussions."
 PDM acknowledges support for this work provided bx NASA through ANAFE Fellowship erant number PES-10005. awarded by the ANAK Science Center. which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NASS-39073.," TDM acknowledges support for this work provided by NASA through AXAF Fellowship grant number PF8-10005 awarded by the AXAF Science Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-39073."
closure phase.,closure phase.
 As the application of these techniques to optical/IR. wavelengths is still a relatively new field. (here is considerable uncertainty. regarding (he measurement. precision that will be achievable.," As the application of these techniques to optical/IR wavelengths is still a relatively new field, there is considerable uncertainty regarding the measurement precision that will be achievable."
 Thus we will define a range of likely precisions (Lor the optimist aud (he pessimist) and evaluate the usefulness of the technique for each case., Thus we will define a range of likely precisions (for the optimist and the pessimist) and evaluate the usefulness of the technique for each case.
 Astrometric measurements of faint sources require the presence of a bright Uy<13) phase reference within an isoplanatie patch. twpically of order 20 arcseconds in the ἐν baa.," Astrometric measurements of faint sources require the presence of a bright $K<13$ ) phase reference within an isoplanatic patch, typically of order 20 arcseconds in the $K$ band."
" The probability of a A<13 star falling within 20"" is roughly (Shao&Colavila1992b).", The probability of a $K<13$ star falling within $20\arcsec$ is roughly \citep{shaocolavita92b}.
. Narrow-angle astrometry at the LOO micro-arcseconcl level has already been demonstrated al PTI. and it is reasonable {ο expect the heck and VLTI to produce the specified precision of 10-30 microarcseconds.," Narrow-angle astrometry at the 100 micro-arcsecond level has already been demonstrated at PTI, and it is reasonable to expect the Keck and VLTI to produce the specified precision of 10-30 microarcseconds."
 However. narrow-angle interferometric astromeltrv is (he least mature of the (11ος techniques. and hence we will consider the utility of the technique for precisions of both LOO and LO µας.," However, narrow-angle interferometric astrometry is the least mature of the three techniques, and hence we will consider the utility of the technique for precisions of both 100 and 10 $\mu$ as."
 Next consider measurement of the visibility amplituce for faint sources., Next consider measurement of the visibility amplitude for faint sources.
 The precision with which |V| can be measured with small apertures (single-ry. where ry is the Fried parameter) ean be as good as0.," The precision with which $|V|$ can be measured with small apertures $r_0$, where $r_0$ is the Fried parameter) can be as good as."
3%.. However. recall that in our estimate for the number of photons per coherence (ime. we assumed an aperture of 811. Since a (vpical Fried parameter [or A band is ry~60 em. adaptive optics (AQ) are required to correct distortions in the wavelront.," However, recall that in our estimate for the number of photons per coherence time, we assumed an aperture of 8m. Since a typical Fried parameter for $K$ band is $r_0\sim60$ cm, adaptive optics (AO) are required to correct distortions in the wavefront."
 Ideally. an AO svstem would completely flatten the wavelront. giving uniform path delays across the aperture.," Ideally, an AO system would completely flatten the wavefront, giving uniform path delays across the aperture."
 This would restore (he point spread Dunction (PSF) to the familiar Airy pattern., This would restore the point spread function (PSF) to the familiar Airy pattern.
 Unfortunately. AO systems are not ideal. and in practice manage to eet ~50% of the light into a tight core. and the remainder in a diffuse halo.," Unfortunately, AO systems are not ideal, and in practice manage to get $\sim50\%$ of the light into a tight core, and the remainder in a diffuse halo."
 This Strehl ratio sels the upper limit to (he (sequared) visibilities an interferometer can achieve. unless a spatial filler such as a single-mode optical fiber is used.," This Strehl ratio sets the upper limit to the (squared) visibilities an interferometer can achieve, unless a spatial filter such as a single-mode optical fiber is used."
 However. such a spatial filler increases ihe maxinum achievable visibili amplitude at the expense of losing lisht. (max. coupling efficiency ~78%. Shaklan&BRoddier (1988))) and introducing visibility biases which depend on the Strehl ratio (Guvon2002).," However, such a spatial filter increases the maximum achievable visibility amplitude at the expense of losing light (max coupling efficiency $\sim78\%$, \citet{shak88}) ) and introducing visibility biases which depend on the Strehl ratio \citep{guyon02}."
. In principle. these effects may be calibrated by observing point sources and monitoring the AO Strehl ratio in real time.," In principle, these effects may be calibrated by observing point sources and monitoring the AO Strehl ratio in real time."
 In practice. such calibration is difficult ancl as vel unproven.," In practice, such calibration is difficult and as yet unproven."
 At present. it is not known how well visibility amplitudes mav be calibrated for large. AO-corrected apertures. but it may potentially be on the order of5%.," At present, it is not known how well visibility amplitudes may be calibrated for large, AO-corrected apertures, but it may potentially be on the order of."
. In summary. the expected precision in visibility amplitude (|V]|) for the Neck and VLTI will likely be between 0.5 and55: we will consider the resulting measurement SNR for both cases.," In summary, the expected precision in visibility amplitude $|V|$ ) for the Keck and VLTI will likely be between 0.5 and; we will consider the resulting measurement SNR for both cases."
 Lastly. we consider closure phase.," Lastly, we consider closure phase."
 As mentioned earlier. closure phase is [free of many of the error sources afflicting visibility amplitude ancl phase.," As mentioned earlier, closure phase is free of many of the error sources afflicting visibility amplitude and phase."
 However. other svstematics can arise which we have not anticipated.," However, other systematics can arise which we have not anticipated."
 The current generation of interferometers have demonstrated closure phase precision at roughly the I-deeree level., The current generation of interferometers have demonstrated closure phase precision at roughly the 1-degree level.
 However. there is no," However, there is no"
Since oy and τω have the samefunctional form (Pi Cry). these two equations are in fact the same thing.,"Since $\psi_1$ and $\psi_2$ have the samefunctional form $P_\ell^m(x)$ ), these two equations are in fact the same thing."
 Iu actual uuuerical procedure. we solve for g; as opposed to c; ," In actual numerical procedure, we solve for $g_i$ as opposed to $\psi_i$ ."
Equation is modified for g; as. HO," Equation is modified for $g_i$ as, =."
"""m Usiug equation(2.2.0).. we find that gy is a polvnonial of order (fJaf)? For cach (i) pair. there can therefore be {ΕΠ} roots satisfying this boundary conditiou. with half of them having jp2 0."," Using equation, we find that $g_1$ is a polynomial of order $(\ell - |m|)$ For each $(\ell, m)$ pair, there can therefore be $(\ell -
|m|)$ roots satisfying this boundary condition, with half of them having $\mu > 0$ ."
" Another wav of plaasing this is that. for cach eigeufuuction (o6)=PI""sPII sy). there are (6ln]) eieenfrequencies."," Another way of phrasing this is that, for each eigenfunction $\psi(x_1,x_2) = P_\ell^{|m|}(x_1)\,
P_\ell^{|m|}(x_2)$ ), there are $(\ell - |m|)$ eigenfrequencies."
 Roushlvhalf of these are progracde modes (202» 0) and the other half retrograde modes (in cU), Roughlyhalf of these are prograde modes $m > 0$ ) and the other half retrograde modes $m < 0$ ).
 Table 1. preseuts some example ceieeufrequencies for iuertialiuodes in a uniforiu density sphere., Table \ref{tab:constantrho} presents some example eigenfrequencies for inertial-modes in a uniform density sphere.
 Our results agree with those presented i Tables 3-1 o£ LF aud.Table lof(1999)., Our results agree with those presented in Tables 3-4 of LF andTable 1 of.
. Iu the previous section. we have introduced various eieenvalues (I. A. ( and ii). as well as the cigcutrequency μι," In the previous section, we have introduced various eigenvalues $K$, $\lambda$, $\ell$ and $m$ ), as well as the eigenfrequency $\mu$."
 What is the geometrical meaning of these eigenvalues and what isthe dispersion relation for inertialanodes?, What is the geometrical meaning of these eigenvalues and what isthe dispersion relation for inertial-modes?
 ere. we present a derivation that answers these questions.," Here, we present a derivation that answers these questions."
 We first convert the imdepeudent variable from ur; to O;=costa; (to be differentiated with the spherical augle 0)so that equation beconies. Lp»11)——ον | X g; 0.," We first convert the independent variable from $x_i$ to $\Theta_i =
\cos^{-1} x_i$ (to be differentiated with the spherical angle ${\theta}$ )so that equation becomes, + (2 |m|+1) + ^2 g_i = 0."
 Adopting a WKD approach. we express 4$;Xexp{hodO;) with the wave vector koη|i; where the real part ky~OCA)z91 aud the imaginary part hy~OL).," Adopting a WKB approach, we express $g_i \propto \exp(i\int k_\Theta d\Theta_i)$ with the wave vector $k_\Theta = k_R+ i k_I$ where the real part $k_R
\sim {\cal O}(\lambda) \gg 1$ and the imaginary part $k_I \sim {\cal O}(1)$."
" Substituting this iuto equation(31).. and equating terms of comparable magnitudes. we fud Together with equation(2.2.0).. this gives rise to the following approximate solution forc; in the WKB τοσο, ∣∖ dn "," Substituting this into equation, and equating terms of comparable magnitudes, we find Together with equation, this gives rise to the following approximate solution for$\psi_i$ in the WKB regime, _i = ^m (x_i) _i + )."
So c; Is an oscillating function with a roughly coustaut envelope., So $\psi_i$ is an oscillating function with a roughly constant envelope.
 Moreover. à=Απ.π for even-parity modes (eq. [2.2.0]:," Moreover, $\alpha = -\lambda\pi/2 + \pi$ for even-parity modes (eq. \ref{eq:boundary2}] ]);"
 while a—Az/2|x/2 for odd-parity modes (eq. [2.2.0]., while $\alpha= -\lambda\pi/2 + \pi/2$ for odd-parity modes (eq. \ref{eq:boundary3}] ]).
 Nodes of c;=0 are roughly evenly spaced in Oc[0.2/2] space. and the value of A corresponds with twice the total απο of nodes (nore below).," Nodes of $\psi_i = 0$ are roughly evenly spaced in $\Theta \in [0,\pi/2]$ space, and the value of $\lambda$ corresponds with twice the total number of nodes (more below)."
 To obtain the dispersion relation. we iusert the above approsiate solution iuto the boundary condition at the surface (eq. [29]|)," To obtain the dispersion relation, we insert the above approximate solution into the boundary condition at the surface (eq. \ref{eq:boundary1}] ])"
 where Oy=coxtyr., where $\Theta_0 = \cos^{-1} \pm \mu$.
 This yields. yyy) -," This yields, ) + ) 0."
" For À»I. this can be approximated by Theta, ΠΠ pe where s ds an integer."," For $\lambda \gg 1$, this can be approximated by _0 + ) +, where $n$ is an integer."
 We retain ouly the positive sigu as) can be positive or negative., We retain only the positive sign as $m$ can be positive or negative.
" For eveu-parity modes. the dispersion relation for the eigeufrequeucy rus as while for odd-parity. it is If we define the number of nodes in the 0.Oy] range to be my. and the umuberin the Ου€ range tobe na. they satisfy ""απ= AOy. HomEÀ(xj2 Quy)."," For even-parity modes, the dispersion relation for the eigenfrequency runs as while for odd-parity, it is If we define the number of nodes in the $\Theta_1 \in [0,\Theta_0]$ range to be $n_1$ , and the numberin the $\Theta_2 \in
[\Theta_0,\pi/2]$ range tobe $n_2$ , they satisfy $n_1
\pi \approx \lambda
\Theta_0$ , $n_2 \pi \approx \lambda (\pi/2 -
\Theta_0)$ ."
 So Azμη| no) as werave 1nentioned earlier.," So $\lambda \approx 2 (n_1 + n_2)$ , as wehave mentioned earlier."
 Moreover. equation vields HomonaLl|onmfwΑπ for even modes. aud à3- for. odd modes.," Moreover, equation yields $n
\approx n_2 - 1 + m/\sqrt{1-\mu^2}\lambda\pi$ for even modes, and $n \approx n_2 - 1/2 + m/\sqrt{1-\mu^2}\lambda\pi$ for odd modes."
 Substitutingη+ hese estimates iuto equations (37).. we find the ‘ollowing simplified dispersion relation.," Substituting these estimates into equations , we find the following simplified dispersion relation, ="
Iu other words. the optically thick enission is dominated by cutting elements ou the near aud far sides of the black hole. for which the Extraordinary wave has a polarization direction parallel to the refercuce axis.,"In other words, the optically thick emission is dominated by emitting elements on the near and far sides of the black hole, for which the Extraordinary wave has a polarization direction parallel to the reference axis."
 In contrast. the dominant contribution iu the thin region comes from the blue shifted cmitter to the side of the black hole. where the Extraordinary wave has a polarization direction mostly perpeucdicular to this axis.," In contrast, the dominant contribution in the thin region comes from the blue shifted emitter to the side of the black hole, where the Extraordinary wave has a polarization direction mostly perpendicular to this axis."
 The sharp decrease in polarization at still higher frequencies is due to the diluting effects of Comptonization emission which beeius to dominate over Svuclirotron cussion at that point., The sharp decrease in polarization at still higher frequencies is due to the diluting effects of Comptonization emission which begins to dominate over Synchrotron emission at that point.
 The inclination anele dependence of the fractional polarization associated with emission bv the Neplerian portion of the inflow is shown in Fieure 2., The inclination angle dependence of the fractional polarization associated with emission by the Keplerian portion of the inflow is shown in Figure 2.
 Several issues remain to be investigated., Several issues remain to be investigated.
 Au important result of our analysis is that only modes accretion rates appear to be consistent with the polarization characteristics of Ser À* at nuu aud sub-uua wavelengths., An important result of our analysis is that only modest accretion rates appear to be consistent with the polarization characteristics of Sgr A* at mm and sub-mm wavelengths.
 The cmitting regiou is compactevidently uo lavecr than a handful of Sclavarzsclild radi., The emitting region is compact—evidently no larger than a handful of Schwarzschild radii.
 Yet hydrodyuaimnical simulatious (Coker Molia 1997) suggest that the rate at which plasma is captured at larger radii (of order 104r& or so} is several orders of magnitude higher., Yet hydrodynamical simulations (Coker Melia 1997) suggest that the rate at which plasma is captured at larger radii (of order $10^4\;r_S$ or so) is several orders of magnitude higher.
 If oux modeling is correct. this would seenài to suggest that AL is variable. perhaps due to a gradual loss of mass with decreasing radius (sec. e.g.. Dlandford Degelinan 1999).," If our modeling is correct, this would seem to suggest that $\dot M$ is variable, perhaps due to a gradual loss of mass with decreasing radius (see, e.g., Blandford Begelman 1999)."
 It is essential to selt-consisteutly match the conditious within- the Ixepleriau region of the flow with the quasi-phlierieal infall further out., It is essential to self-consistently match the conditions within the Keplerian region of the flow with the quasi-spherical infall further out.
 These caleulatious are currently under wax. and the results will be reported clsewhere.," These calculations are currently under way, and the results will be reported elsewhere."
 In addition. if the source is indeed the counterpart to Ser A* (Daganoff et al 2000). then the spectrum shown in Figure 1 sugeests a correlated variability between the sib-uun aud N-rav fluxes. which can be tested with the next round of coordinated observations.," In addition, if the source is indeed the counterpart to Sgr A* (Baganoff et al 2000), then the spectrum shown in Figure 1 suggests a correlated variability between the sub-mm and X-ray fluxes, which can be tested with the next round of coordinated observations."
 We are very erateful to Marco Fatuzzo for helpful discussions., We are very grateful to Marco Fatuzzo for helpful discussions.
 This work was supported by a Sir Thomas Lyle Fellowship aud a Mieguuyahli Fellowship for distinguished overseas visitors at the University of Melbourne. and by NASA erants NAC5-8239 and NÀCO-9205.," This work was supported by a Sir Thomas Lyle Fellowship and a Miegunyah Fellowship for distinguished overseas visitors at the University of Melbourne, and by NASA grants NAG5-8239 and NAG5-9205."
(Jonkeretal.2004.2007:Tomsick2005:Wijnancsetal.2005:Lleinke 2009).,"\citep{jonker2004_saxj1810,jonker07,tomsick2005,wijnands2005,heinke2009}."
. These neutron stars must oe cold. suggesting that their interior may be elliciently cooled through neutrino emissions (see also Sec. 22)).," These neutron stars must be cold, suggesting that their interior may be efficiently cooled through neutrino emissions (see also Sec. \ref{sec:thermal}) )."
 Such neutron stars are expected to be relatively massive. because arger central densities are thought to lead to a higher rate of neutrino cooling (Lattimer&Prakash2004:YakovlevPothick 2004).," Such neutron stars are expected to be relatively massive, because larger central densities are thought to lead to a higher rate of neutrino cooling \citep[][]{lattimer2004,yakovlev2004}."
. In this work. we discuss three additional oobservations of Terzan 5. carried out in 2009 and 2011. to urther investigate the quiescent properties of245.," In this work, we discuss three additional observations of Terzan 5, carried out in 2009 and 2011, to further investigate the quiescent properties of."
. We use four oobservations of Terzan 5. spread: between 2003 ancl 2011. to study any. possible spectral and temporal variability in the quiescent emission. of248.," We use four observations of Terzan 5, spread between 2003 and 2011, to study any possible spectral and temporal variability in the quiescent emission of."
.. Details of the individual observations can be found in Table 2. anc references therein., Details of the individual observations can be found in Table \ref{tab:obs} and references therein.
 All four were carried out in the faint data mode. with the nominal frame time of 3.2 s and the targe positioned on the $3 chip.," All four were carried out in the faint data mode, with the nominal frame time of 3.2 s and the target positioned on the S3 chip."
 Data reduction was carried ou using the tools (v. 4.3) and following standard: aanalvsis Phe 2003 and 2009data were reprocesse using the task EVENTS to benefit. [rom the most recent. calibration., Data reduction was carried out using the tools (v. 4.3) and following standard analysis The 2003 and 2009data were reprocessed using the task $\_$ $\_$ EVENTS to benefit from the most recent calibration.
 IEZpisodes of high backgrotunc were removed from the 2003 data. which resulted in a ne exposure time of 31.2 ks.," Episodes of high background were removed from the 2003 data, which resulted in a net exposure time of 31.2 ks."
 No background fares were presen during the 2009 and 2011 observations. so the full exposure time was used in the analysis.," No background flares were present during the 2009 and 2011 observations, so the full exposure time was used in the analysis."
 iis clearly detected in the dense core of the cluster in 2007Ü and 2009 (Eig. 2.. ," is clearly detected in the dense core of the cluster in 2003 and 2009 (Fig. \ref{fig:ds9}, ,"
loft)., left).
" We used a circular region with a 1.5η radius to extract source events. and one with a radius of 40"" positioned on à source-Iree part of the CCD ~L4 west of the cluster core as a background. reference."," We used a circular region with a $1.5''$ radius to extract source events, and one with a radius of $40''$ positioned on a source-free part of the CCD $\sim1.4'$ west of the cluster core as a background reference."
 We extracted count rates and lishteurves using the tool DALENPRACT. while the meta-task SPECENTRACT was used to obtain spectra and to generate the ancillary response files (arf) ancl redistribution matrix files (rmf).," We extracted count rates and lightcurves using the tool DMEXTRACT, while the meta-task SPECEXTRACT was used to obtain spectra and to generate the ancillary response files (arf) and redistribution matrix files (rmf)."
 Phe spectral cata was grouped into bins with a minimum of 15 photons using the tool GRPPILA and. fitted using the software package (v. 12.7)., The spectral data was grouped into bins with a minimum of 15 photons using the tool GRPPHA and fitted using the software package (v. 12.7).
 To take into account the interstellar neutral hyclrogen absorption along the line of sight. we include the PLLABS model in all our spectral fits using the default. abundances and. cross-sections.," To take into account the interstellar neutral hydrogen absorption along the line of sight, we include the PHABS model in all our spectral fits using the default abundances and cross-sections."
 Throughout this work we assume a distance of D=5.5 kpe towards wwhen converting N-ray fluxes into luminosities., Throughout this work we assume a distance of $D=5.5$ kpc towards when converting X-ray fluxes into luminosities.
 ALL quoted errors correspond to 90% confidence levels., All quoted errors correspond to $90\%$ confidence levels.
 Whereas wwas one of the brightest sources in Terzan 5 during the 2003 and 2009 observations. the source had. considerably faced in 2011 (Fig. 2)).," Whereas was one of the brightest sources in Terzan 5 during the 2003 and 2009 observations, the source had considerably faded in 2011 (Fig. \ref{fig:ds9}) )."
 Although the source is not clearly visible in the 2011 data sets. there appears to be an excess of photons present at the source position.," Although the source is not clearly visible in the 2011 data sets, there appears to be an excess of photons present at the source position."
 To verify this. we emploved the wavelet detection algorithm. PNDETECT (Damianietal.1997a.b).. which has been found to be particularly effective in detecting faint sources located close to brighter objects and henee it is a very useful tool for globular clusters (c.g..Heinkeetal. 2006)...," To verify this, we employed the wavelet detection algorithm PWDETECT \citep{damiani1997_1,damiani1997_2}, which has been found to be particularly effective in detecting faint sources located close to brighter objects and hence it is a very useful tool for globular clusters \citep[e.g.,][]{heinke2006_terzan5}. ."
" We performedhiandard PWDIZEEC"" runs on the ACIS-S3 chip with wavelet sizes varving E.from 0.5”2.07! ""phis. exercise", We performedstandard PWDETECT runs on the ACIS-S3 chip with wavelet sizes varying from $0.5''-2.0''$ This exercise
the SN Ia rate may be less in LSBs than HISDs. but not signilicantlv less in ellipticals aud early-type spirals.,"the SN Ia rate may be less in LSBs than HSBs, but not significantly less in ellipticals and early-type spirals."
 Though (he assumption that the SN Ia rate in LSBs is the same as that in IISBs is reasonable. it is still an assumption. aud it may be modified as supernovae in LSBs are studied further.," Though the assumption that the SN Ia rate in LSBs is the same as that in HSBs is reasonable, it is still an assumption, and it may be modified as supernovae in LSBs are studied further."
 Cappellaroetal.(1999). determine that the SN Ia rate in ISB galaxies is 0.1840.05 SNe centurv.HO.ILus. κο a lower limit on the SN la rate of >0.13 SNe ! 10.IL. is adopted.," \citet{cap99} determine that the SN Ia rate in HSB galaxies is $0.18 \pm 0.05$ SNe $^{-1} 10^{-10} L_{B,\odot}$, so a lower limit on the SN Ia rate of $\ge 0.13$ SNe $^{-1}$ $10^{-10} L_{B,\odot}$ is adopted."
 Dividing (he upper limits on the SN Ia rate per volume in LSB galaxies bv the lower limit on the SN Ia rate per mass (Cappellaroetal.1999). vields upper limits on the contribution of LSBs to the optical Iuminosityv density of the local Universe., Dividing the upper limits on the SN Ia rate per volume in LSB galaxies by the lower limit on the SN Ia rate per mass \citep{cap99} yields upper limits on the contribution of LSBs to the optical luminosity density of the local Universe.
 We have calculated a mean LSB barvonic mass-to-light ratio of 2.20 (solar units) from the results of deBloketal.(1996).. Matthewsetal. (2001).. MeGaugh&deBlok(1997 ).. MonnierRagaigneetal. (2003a.b).. O'Neil.Bothun.&Schombert(2000).. and Zavala (2003).," We have calculated a mean LSB baryonic mass-to-light ratio of 2.20 (solar units) from the results of \citet{deb96}, \citet{mat01}, \citet{mcg97}, , \citet{mon03a,mon03b}, , \citet{one00}, and \citet{zav03}."
. From the data of deBloketal.(1996).. MeGaugh&deBlok (1998).. aud. (2003).. we calculate a mean LSB dynamical mass-to-light ratio of 20.21.," From the data of \citet{deb96}, \citet{mcg98}, and \citet{zav03}, we calculate a mean LSB dynamical mass-to-light ratio of 20.21."
 Note that the dvnamical ratio is less certain than the barvonic ratio because the number of LSBs with ealeulated harvonic mass-to-light ratios is much greater than the number of LSBs with calculated dynamical mass-to-light ratios., Note that the dynamical ratio is less certain than the baryonic ratio because the number of LSBs with calculated baryonic mass-to-light ratios is much greater than the number of LSBs with calculated dynamical mass-to-light ratios.
" We have used these (wpical mass-to-light ratios to set upper limits on the contribution of LSBs to OQ, and O,,.", We have used these typical mass-to-light ratios to set upper limits on the contribution of LSBs to $\Omega_b$ and $\Omega_m$ .
 The results of the caleulations are listed in Table 1.., The results of the calculations are listed in Table \ref{tab:limits}.
 The optical luminosity density. £g. of ΗΡΙ galaxies is (1.3520.14)x10Lg.Mpe* (Fukueita&Peebles2004).," The optical luminosity density, $\cal L_B$, of HSB galaxies is $(1.35 \pm 0.14) \times
10^8 \, L_{B,\odot} \, \hbox{Mpc}^{-3}$ \citep{fuk04}."
.! The le and confidence upper limits on the optical luminosity density of LSBs are slightly above (he HSB value. so LSBs and displaced stars contribute no more (han an amount comparable to LSBs to theoptical luminosity density of the local Universe.," The $1\sigma$ and confidence upper limits on the optical luminosity density of LSBs are slightly above the HSB value, so LSBs and displaced stars contribute no more than an amount comparable to HSBs to theoptical luminosity density of the local Universe."
" We set upper limits on the contribution of LSBs to the mean Q, aud Q,,, of the local Universe (Table 1)).", We set upper limits on the contribution of LSBs to the mean $\Omega_b$ and $\Omega_m$ of the local Universe (Table \ref{tab:limits}) ).
" The upper limit on ©, lor LSBs indicates that LSBs cannot increase the currently observed mean value of ος in the Universe bx more than 9 percent. because Q,=0.044+0.004 (Bennettetal.2003).", The upper limit on $\Omega_b$ for LSBs indicates that LSBs cannot increase the currently observed mean value of $\Omega_b$ in the Universe by more than 9 percent because $\Omega_b = 0.044 \pm 0.004$ \citep{ben03}.
. Therefore. LSBs are not significant. barvonic repositories in (he local Universe.," Therefore, LSBs are not significant baryonic repositories in the local Universe."
" LSBs cannot increase (he value of the mean ©,, in the Universe by more (han 13 percent because O,,=0.27+0.04 (Bennettetal. 2003)..", LSBs cannot increase the value of the mean $\Omega_m$ in the Universe by more than 13 percent because $\Omega_m = 0.27 \pm 0.04$ \citep{ben03}. .
 Thus LSBs cannot account for a significant component ofthetotal dark matter in (heUniverse., Thus LSBs cannot account for a significant component ofthetotal dark matter in theUniverse.
peaks in the stable region. which reflects the characteristics 1) aud d) of the flow pattern described above.,"peaks in the stable region, which reflects the characteristics i) and ii) of the flow pattern described above."
 Figure 8bb shows radial velocities of particles predicted by equation (13)) using simulated eas velocity alter the establishment of the euasi-steady flow (/O=10)., Figure \ref{fig:s40-dust}b b shows radial velocities of particles predicted by equation \ref{eq:turb}) ) using simulated gas velocity after the establishment of the quasi-steady flow $t \Omega = 70$ ).
 The dots concentrated ab ο)~2.0 in the figure are actual velocities of all the simulated particles., The dots concentrated at $x/H \simeq 2.0$ in the figure are actual velocities of all the simulated particles.
 The actual dala of the particle velocities are well reproduced. by equation (13)). since most ol the dala points are located between the predicted minimum and maxinum values.," The actual data of the particle velocities are well reproduced by equation \ref{eq:turb}) ), since most of the data points are located between the predicted minimum and maximum values."
 We further decompose vy into ep anc ος according to equation (13))., We further decompose $v_{\rm d}$ into $v_{\rm f}$ and $v_{\rm t}$ according to equation \ref{eq:turb}) ).
 Figure δος presents the radial distribution of (he two components., Figure \ref{fig:s40-dust}c c presents the radial distribution of the two components.
 The maximum and minimum values lor each component are plotted., The maximum and minimum values for each component are plotted.
 It is shown that the effect of the remnant turbulence is small aud overall features of particle radial migration is represented by vy., It is shown that the effect of the remnant turbulence is small and overall features of particle radial migration is represented by $v_{\rm f}$ .
 Due to the non-smooth pressure and velocity distributions. min|e4]eπμ> Oat —3.5Z5vc/IHIS—2.5 and r/H~—2.1. in addition to 0.0Srfl2.0.," Due to the non-smooth pressure and velocity distributions, ${\rm min}[v_{\rm d}] \simeq {\rm min}[v_{\rm f}]>0$ at $-3.5 \lesssim x/H \lesssim -2.5$ and $x/H \sim -2.1$, in addition to $0.0 \lesssim x/H \lesssim 2.0$."
 It means that all particles have positive radial velocity there and accumulate near the outer edges of these regions., It means that all particles have positive radial velocity there and accumulate near the outer edges of these regions.
 On the other hand. max[e4]2πανκ>0 is observed in some areas. [or example. around 11~3.0.," On the other hand, ${\rm max}[v_{\rm d}] \simeq {\rm max}[v_{\rm f}]>0$ is observed in some areas, for example, around $x/H \simeq 3.0$."
 A small number of particles are stalled there., A small number of particles are stalled there.
 The peaks of particle density at. /if<0 would not grow because (here are few particles which can accumulate from outside (Figure 8aa)., The peaks of particle density at $x/H<0$ would not grow because there are few particles which can accumulate from outside (Figure \ref{fig:s40-dust}a a).
 The stagnant areas cannot stall all the particles which migrate inward due to minjea]<0 (Figure 8bb)., The stagnant areas cannot stall all the particles which migrate inward due to ${\rm min}[v_{\rm d}]<0$ (Figure \ref{fig:s40-dust}b b).
 Therefore the peak at ΠΠ~2.0 ds the most prominent concentration zone., Therefore the peak at $x/H \simeq 2.0$ is the most prominent concentration zone.
 Figure Yaa shows the (ime variation of the number of particles in the clensest erid in the whole region and al (he center of the traced clump., Figure \ref{fig:s40-clump}a a shows the time variation of the number of particles in the densest grid in the whole region and at the center of the traced clump.
 The agreement between the two inclicates that the same clump is growing., The agreement between the two indicates that the same clump is growing.
"Figure 9bb is the velocity dispersion of particles within the chunp at /=/,(—55/0).",Figure \ref{fig:s40-clump}b b is the velocity dispersion of particles within the clump at $t=t_c (=55/\Omega)$.
 The velocity dispersion is small but is larger than in model-j., The velocity dispersion is small but is larger than in $\eta$.
 While the turbulence does not affect the location of the dust concentration. it does influence the degree of particle accumulation.," While the turbulence does not affect the location of the dust concentration, it does influence the degree of particle accumulation."
 When the azimuthal box size is long. shear velocity becomes supersonic near the edge and it may cause artificial densitv dip (Johnsonetal.2008:Johansen2009) unless a hieh-order scheme is applied for Keplerian advection term.," When the azimuthal box size is long, shear velocity becomes supersonic near the edge and it may cause artificial density dip \citep{john08, joh09} unless a high-order scheme is applied for Keplerian advection term."
 To confirm that the density dip [found in our simulations is not caused by such a iunerical error. we also carried out a run in which the unstable region is put in the side areas ancl (he stable region is at the center.," To confirm that the density dip found in our simulations is not caused by such a numerical error, we also carried out a run in which the unstable region is put in the side areas and the stable region is at the center."
 The other parameters are (he same as mocdel-s40., The other parameters are the same as model-s40.
 We found that the density dip is formed al the side area and no artificial density dip is created in the central area. which means that our numerical scheme using CIP method with thestaggard mesh and the short timestep (dl=10. ἐν Courant condition is set to be < 0.5)- is sufficiently high-order scheme.," We found that the density dip is formed at the side area and no artificial density dip is created in the central area, which means that our numerical scheme using CIP method with thestaggard mesh and the short timestep $dt=10^{-4}$ ; Courant condition is set to be $<0.5$ ) is sufficiently high-order scheme."
"As we have seen in Sec.??,, during the later stages of the pre main sequence evolution captured DM can become an important energy source.","As we have seen in \ref{dmcapture}, during the later stages of the pre main sequence evolution captured DM can become an important energy source."
 The luminosity due to DM capture ts where falyr and the factor of 2 in Eq.(??)) reflects the fact that the energy per annihilation is twice the WIMP mass., The luminosity due to DM capture is where dV and the factor of $2$ in \ref{caplum}) ) reflects the fact that the energy per annihilation is twice the WIMP mass.
" In all simulations we will consider the case of ""minimal capture"", which corresponds to equal contribution to the luminosity from capture and nuclear fusion when the star reaches the main sequence."," In all simulations we will consider the case of “minimal capture”, which corresponds to equal contribution to the luminosity from capture and nuclear fusion when the star reaches the main sequence."
" On the whole. for all the values of boost factor and concentration parameter we have considered in this paper, the results are roughly the same: the final DS is roughly ~1000... ον10*L.., and lives ~LO’ yrs."," On the whole, for all the values of boost factor and concentration parameter we have considered in this paper, the results are roughly the same: the final DS is roughly $\sim 1000 \msun$ , $\sim 10^7 L_\odot$, and lives $\sim 10^6$ yrs."
" Thus if the e exeess seen in PAMELA is due to WIMP annihilation, the required leptophilic boosted cross section is certainly compatible with the DS picture."," Thus if the $e^+$ excess seen in PAMELA is due to WIMP annihilation, the required leptophilic boosted cross section is certainly compatible with the DS picture."
 However there are interesting differences between models which we will discuss., However there are interesting differences between models which we will discuss.
" Other than in the subsection immediately following this one, we will consider four WIMP models."," Other than in the subsection immediately following this one, we will consider four WIMP models."
" As motivated below, we will focus on one boosted model denoted by AH4 with the following set of parameters: D— 1500.70,=2.25 TeV and ¢=2.5."," As motivated below, we will focus on one boosted model denoted by AH4 with the following set of parameters: $B=1500$, $m_{\chi}=2.35$ TeV and $c=3.5$."
" As our unboosted models. we will take 100 GeV WIMPs with the canonical cross section of 3x10."" em""/s, and we will consider three values of the concentration parameter, c = (2. 3.5. 5)."," As our unboosted models, we will take 100 GeV WIMPs with the canonical cross section of $3 \times 10^{-26}$ $^3$ /s, and we will consider three values of the concentration parameter, c = (2, 3.5, 5)."
" For comparison. the ""canonical case” considered in ? was the unboosted 100 GeV case with c=2."," For comparison, the ""canonical case"" considered in \citet{DSnl} was the unboosted 100 GeV case with c=2."
" The relative boost factor between the AH4 model and the unboosted models is best described as follows: since Eq.(1) tells us that DM heating scales as (ov)/M,. one can see thatthe AH4 is exactly equivalent to a 100 GeV WIMP"," The relative boost factor between the AH4 model and the unboosted models is best described as follows: since Eq.(1) tells us that DM heating scales as $\sv/M_\chi$, one can see thatthe AH4 is exactly equivalent to a 100 GeV WIMP"
"distribution of normal and compact high-z ETGs, in Fig.","distribution of normal and compact $z$ ETGs, in Fig."
 8 we report colour gradient as a function of the galaxy degree of compactness (see Sect., \ref{gradcomp} we report colour gradient as a function of the galaxy degree of compactness (see Sect.
 2)., 2).
" Circles represent compact galaxies, while triangles are normal ones."," Circles represent compact galaxies, while triangles are normal ones."
 The open points represent galaxies for which Sersic indices are locked in both bands., The open points represent galaxies for which Sersic indices are locked in both bands.
" Since simulations leave a room of space to uncertainties into the estimates of the colour gradients in these systems, we prefer not to take into account these points."," Since simulations leave a room of space to uncertainties into the estimates of the colour gradients in these systems, we prefer not to take into account these points."
 We observe that the present galaxy sample seems to show a mild correlation between the degree of compactness and colour gradient., We observe that the present galaxy sample seems to show a mild correlation between the degree of compactness and colour gradient.
" Indeed, compact galaxies seem to preferentially show a bluer core than the outer regions, while moving towards normal galaxies stellar populations become redder in the center as observed in the local Universe."," Indeed, compact galaxies seem to preferentially show a bluer core than the outer regions, while moving towards normal galaxies stellar populations become redder in the center as observed in the local Universe."
" In fact, only one compact galaxy out of 6 (15%) has a negative colour gradient (but compatible within the error with a positive value) against the 45% in the case of normal galaxies."," In fact, only one compact galaxy out of 6 $\%$ ) has a negative colour gradient (but compatible within the error with a positive value) against the $\%$ in the case of normal galaxies."
 This difference is detectable in the mean colour gradient value of the two samples., This difference is detectable in the mean colour gradient value of the two samples.
" Indeed, while the mean colour gradient of compact galaxies is <Vuv-—u>comp 0.22+0.28, normal galaxies have a mean value of <Vuv—u>norm 0.04+0.30."," Indeed, while the mean colour gradient of compact galaxies is $<$$\nabla$$_{UV-U}$$>_{comp}$ $\pm$ 0.28, normal galaxies have a mean value of $<$$\nabla$$_{UV-U}$$>_{norm}$ $\pm$ 0.30."
" Even if the two mean estimates seem to point towards a different nature of the two samples, it is to note that they are consistent within the errors and that the KS probability for the two samples of being extracted from different populations is not significant (~26%)."," Even if the two mean estimates seem to point towards a different nature of the two samples, it is to note that they are consistent within the errors and that the KS probability for the two samples of being extracted from different populations is not significant $\sim$ $\%$ )."
" At the same time, the Spearman’s rank test with a coefficient p ==--0.34 shows that the correlation between the two quantities has a probability of ~ 85% to not be observed by chance, but even in this case, the poor statistics of our sample prevent us to reach a firm conclusion."," At the same time, the Spearman's rank test with a coefficient $\rho$ -0.34 shows that the correlation between the two quantities has a probability of $\sim$ $\%$ to not be observed by chance, but even in this case, the poor statistics of our sample prevent us to reach a firm conclusion."
" In fact, the trend slightly hinted in Fig."," In fact, the trend slightly hinted in Fig."
 8 can be even due to the incompleteness of our sample., \ref{gradcomp} can be even due to the incompleteness of our sample.
" Unfortunately, we are not able to discriminate how and in which direction (if any) the selection criteria of our sample can arrange the distribution of the points in Fig. 8,,"," Unfortunately, we are not able to discriminate how and in which direction (if any) the selection criteria of our sample can arrange the distribution of the points in Fig. \ref{gradcomp},"
 and hence the effect they could have on the supposed connection., and hence the effect they could have on the supposed connection.
" Thus, although the quality of our sample allow us to investigate the internal colour distribution of both compact and normal high-z ETGs, we cannot be conclusive about the suggested physical connection between the degree of compactness and radial colour variations."," Thus, although the quality of our sample allow us to investigate the internal colour distribution of both compact and normal $z$ ETGs, we cannot be conclusive about the suggested physical connection between the degree of compactness and radial colour variations."
 We analysed restframe colour gradients in a sample of 20 ETGs at (UV-U)0.9<Zspec «1.92 comprising both normal and compact galaxies., We analysed $_{rest frame}$ colour gradients in a sample of 20 ETGs at $<z_{spec}<$ 1.92 comprising both normal and compact galaxies.
" Despite the short wavelength baseline covered, we detect effective radial colour variations in ~ 5096 of the galaxies of our sample, while the remaining show a gradient comparable with zero."," Despite the short wavelength baseline covered, we detect effective radial colour variations in $\sim$ $\%$ of the galaxies of our sample, while the remaining show a gradient comparable with zero."
" In particular, we detect significant radial colour variations in 10 out of the 20 galaxies of our sample: five galaxies exhibit a reddening towards the internal regions (negative colour gradient), as generally observed in the local Universe, and five galaxies show stellar populations bluer in the central regions than in the outskirts, which reflects in positive colour gradients."," In particular, we detect significant radial colour variations in 10 out of the 20 galaxies of our sample: five galaxies exhibit a reddening towards the internal regions (negative colour gradient), as generally observed in the local Universe, and five galaxies show stellar populations bluer in the central regions than in the outskirts, which reflects in positive colour gradients."
 This result supports the previous findings on colour gradients at intermediate and high redshifts (e.g.??7) which already revealed the presence of blue core ETGs which are not common in the local Universe.," This result supports the previous findings on colour gradients at intermediate and high redshifts \citep[e.g.][]{menanteau01,moth02,ferreras09} which already revealed the presence of blue core ETGs which are not common in the local Universe."
" Taking advantage of our composite sample including both compact and normal ETGs, we investigate possible systematics in the stellar content of these physically different systems."," Taking advantage of our composite sample including both compact and normal ETGs, we investigate possible systematics in the stellar content of these physically different systems."
" Actually, the distribution of our data in the plane defined by colour gradients and compactness is suggestive of a mild correlation whereby the colour gradient decreases from positive to negative values moving from compact towards normal galaxies."," Actually, the distribution of our data in the plane defined by colour gradients and compactness is suggestive of a mild correlation whereby the colour gradient decreases from positive to negative values moving from compact towards normal galaxies."
" This trend reflects in the mean values of colour gradients of compact and normal galaxies, 0.22+0.28 and 0.04+0.30 respectively, even if the difference is not statistically significant to point toward two strictly different mass assembly histories."," This trend reflects in the mean values of colour gradients of compact and normal galaxies, $\pm$ 0.28 and $\pm$ 0.30 respectively, even if the difference is not statistically significant to point toward two strictly different mass assembly histories."
" Concurrently, the Spearman’s test assess that the correlation between ETGs compactness and colour gradient could be present with a probability of ~ 85% suggesting a possible physical connection between the two quantities."," Concurrently, the Spearman's test assess that the correlation between ETGs compactness and colour gradient could be present with a probability of $\sim$ $\%$ suggesting a possible physical connection between the two quantities."
" However, we cannot reach a firm conclusion since the poor statistics and the incompleteness could concur to originate a fake correlation."," However, we cannot reach a firm conclusion since the poor statistics and the incompleteness could concur to originate a fake correlation."
" Moreover, it is to note that this trend could be affected by the short wavelength baseline covered which do not sample the emission of the oldest stellar populations."," Moreover, it is to note that this trend could be affected by the short wavelength baseline covered which do not sample the emission of the oldest stellar populations."
" Although the limits of our sample do not allow us to strictly assess or deny this suggested trend, theoretical simulations of the hierarchical assembly of the stellar matter seem to support the presence of blue cores in compact galaxies (?).."," Although the limits of our sample do not allow us to strictly assess or deny this suggested trend, theoretical simulations of the hierarchical assembly of the stellar matter seem to support the presence of blue cores in compact galaxies \citep{hopkins08}."
" Indeed, at higher redshifts, the progenitor galaxies tend to have a greater amount of gas not yet converted in stars."," Indeed, at higher redshifts, the progenitor galaxies tend to have a greater amount of gas not yet converted in stars."
" In a merger event involving a gas rich galaxy, the torsion forces convey this reservoir toward the center giving origin to an intense and short burst of star formation."," In a merger event involving a gas rich galaxy, the torsion forces convey this reservoir toward the center giving origin to an intense and short burst of star formation."
" Concurrently, the pre-existing stars tend to arrange themselves in the external regions leading to the formation of an ETG with a younger (and hence bluer) stellar population in the center than in the outskirts and, consequently, with a positive colour gradient."," Concurrently, the pre-existing stars tend to arrange themselves in the external regions leading to the formation of an ETG with a younger (and hence bluer) stellar population in the center than in the outskirts and, consequently, with a positive colour gradient."
" On the contrary, in later epochs, the progenitor galaxies, having"," On the contrary, in later epochs, the progenitor galaxies, having"
"From the equations 12 and 16 - 20. one obtains the observed EC(bl) energy flux and the observed EC(bl) break energy The Thomson regime condition for the considerecl process reads as The EC(bI) emission is determined by the jet parameters ονκ. ond 7i,5 plus the distance [rom the blazar source r. the unknown magnetic field D. kinematic [actors D and 8 and. finally. by properties of the blazar core [pLyla o.vyo ond Dy.","From the equations 12 and 16 - 20, one obtains the observed EC(bl) energy flux and the observed EC(bl) break energy The Thomson regime condition for the considered process reads as The EC(bl) emission is determined by the jet parameters $[\nu L_{\nu}]_{syn, \,
 br}$ and $\nu_{syn, \, br}$ plus the distance from the blazar source $r$, the unknown magnetic field $B$ , kinematic factors $\Gamma$ and $\theta$ and, finally, by properties of the blazar core $[\nu L_{\nu}]_{bl, \, br}$ ,$\nu_{bl, \, br}$ and $\Gamma_{bl}$."
" For the usually discussed 0c Land 8>45"". typical τι71 and the equipartition value Boy—I. one expects the observed EC(b]) break energy to be Speypry.1—100 TeV in the case of comptonisation of HBL radiation (b/= UBL). and τομ~0.01—1 TeV in the case of LBL-like core emission (bf= LBL)."," For the usually discussed $\delta \sim 1$ and $\theta \geq 45^0$, typical $\nu_{syn, \, 14} \sim 1$ and the equipartition value $B_{-4} \sim 1$, one expects the observed EC(bl) break energy to be $\varepsilon_{ec(HBL), \, br} \sim 1 - 100$ TeV in the case of comptonisation of HBL radiation $bl \equiv HBL$ ), and $\varepsilon_{ec(LBL), \, br} \sim 0.01
 - 1$ TeV in the case of LBL-like core emission $bl \equiv LBL$ )."
" As the svnchrotron break lrequency of the blazar emission is anticorrelated with the blazar svnchrotron Iuminosity. (see section 2.2.2). one can pul upper limits on the observed EC(b1) fluxes [p5,],ppp<10. ο. TeV (taking a rough estimate £151,5,00)<107 erg/s). and [ESappro<LOMDyeere/sem? [or Seetee)10 GeV Owith Leper.(0)<10"" erg/s)."," As the synchrotron break frequency of the blazar emission is anticorrelated with the blazar synchrotron luminosity (see section 2.2.2), one can put upper limits on the observed EC(bl) fluxes $[\nu S_{\nu}]_{ec(HBL), \, br} < 10^{-11} \,
 D_{10}^{-2} \, {\rm erg / s \, cm^2}$ for $\varepsilon_{ec(HBL), \, br} \sim 1$ TeV (taking a rough estimate $L_{HBL, \, max}(0) \leq 10^{46}$ erg/s), and $[\nu S_{\nu}]_{ec(LBL), \, br} < 10^{-10} \, D_{10}^{-2} \, {\rm erg / s \,
 cm^2}$ for $\varepsilon_{ec(LBL), \, br} \sim 10$ GeV (with $L_{LBL, \, max}(0)
 \leq 10^{47}$ erg/s)."
 In addition. the ECUIBL) emission is expected to be significantly decreased due to the KN ellects. contrary to the EC(LDL) radiation.," In addition, the EC(HBL) emission is expected to be significantly decreased due to the KN effects, contrary to the EC(LBL) radiation."
 For the cliseussecl standard jet parameters the latter one could be observed by [rom sources with D<100 Mpc., For the discussed standard jet parameters the latter one could be observed by from sources with $D < 100$ Mpc.
" Similarly to the case of the SSC! process. for a given [rjL,|,pb» the subequipartition magnetic field and/or large jet inclination increase theexpected value of [|jS,uibe- "," Similarly to the case of the SSC process, for a given $[\nu L_{\nu}]_{syn,
 \, br}$ the subequipartition magnetic field and/or large jet inclination increase theexpected value of $[\nu S_{\nu}]_{ec(bl), \, br}$."
From the equations 13 and 16 - 20. one obtains the observed EC(star) energy flux," From the equations 13 and 16 - 20, one obtains the observed EC(star) energy flux"
of finding unbound close pairs increases.,of finding unbound close pairs increases.
 This is taken into account in the cosmological averaged merger timescales (see also 2)., This is taken into account in the cosmological averaged merger timescales (see also ).
 Since. the assumed merger timescale is) the most uncertain quantity im Eq. (9).," Since the assumed merger timescale is the most uncertain quantity in Eq. \ref{mrpar}) ),"
 we compare TM with other recent estimations in the literature., we compare $T_{\rm MM}^{K08}$ with other recent estimations in the literature.
 perform N- body/hydrodynamical simulations of major and minor mergers to study the merger timescales of morphological and close pair approaches., perform $N$ -body/hydrodynamical simulations of major and minor mergers to study the merger timescales of morphological and close pair approaches.
" The principal galaxy in their simulations has log(M,/M.)=10.7. similar to the average mass of our principal galaxies with a close companion. so their major merger timescales. denoted ΤΟMM should be comparable to the previous ieTA."," The principal galaxy in their simulations has $\log\, (M_{\star}/M_{\odot}) = 10.7$, similar to the average mass of our principal galaxies with a close companion, so their major merger timescales, denoted $T_{\rm MM}^{JL10}$ , should be comparable to the previous $T_{\rm MM}^{K08}$."
 We summarize the average values of 777 in Table 7. after correcting with the factor Ομ., We summarize the average values of $T_{\rm MM}^{JL10}$ in Table \ref{tmm} after correcting with the factor $C_{\rm p}$.
" Wefind that Tyrie,<TAN.", Wefind that $T_{\rm MM}^{JL10} < T_{\rm MM}^{K08}$.
" However. the TA include the factor C,,. while the The do not."," However, the $T_{\rm MM}^{K08}$ include the factor $C_{\rm m}$, while the $T_{\rm MM}^{JL10}$ do not."
" Applying to Tene a typical value of C,,=0.6(222?2).. we find thatrates."," Applying to $T_{\rm MM}^{JL10}$ a typical value of $C_{\rm m} = 0.6$, we find that."
 On the other hand. use cosmological simulations to study Ομ and the merger timescale. denoted Teint: ," On the other hand, use cosmological simulations to study $C_{\rm m}$ and the merger timescale, denoted $T_{\rm MM}^{LL10}$."
"They find Tayi”~14 Gyr for logM,/M.)~10.3 galaxies and ry<50/1 kpe (this value includes the factor C4,=0.7 derived from their simulations)."," They find $T_{\rm MM}^{LL10} \sim 1.4$ Gyr for $\log\, (M_{\star}/M_{\odot}) \sim 10.3$ galaxies and $r_{\rm p} \leq 50h^{-1}$ kpc (this value includes the factor $C_{\rm m} = 0.7$ derived from their simulations)."
 This timescale is lower by a factor of two than the one from for this mass. TAN)=2.7 Gyr.," This timescale is lower by a factor of two than the one from for this mass, $T_{\rm MM}^{K08} = 2.7$ Gyr."
 However. assume that the galaxy merger occurs a dynamical friction time after the dark matter halo merger: while do not consider this extra time.," However, assume that the galaxy merger occurs a dynamical friction time after the dark matter halo merger; while do not consider this extra time."
 This fact mitigates the difference between both works. but a more Ple:etailed comparison is needed.," This fact mitigates the difference between both works, but a more detailed comparison is needed."
 In the following we omit the uper index in Z5 for clarity., In the following we omit the super index in $T_{\rm MM}^{\rm K08}$ for clarity.
 The merger rate is an absolute quantity. and should not nmepend on the that we use to infer it.," The merger rate is an absolute quantity, and should not depend on the $r_{\rm p}^{\rm max}$ that we use to infer it."
" Because of this. the —ncrease of the 1)""merger fraction with m (Sect. ??.. "," Because of this, the increase of the merger fraction with $r_{\rm p}^{\rm max}$ (Sect. \ref{ffmu}, ,"
Fig. 4) , Fig. \ref{brpfig}) )
must be compensated with the increase in Zap., must be compensated with the increase in $T_{\rm MM}$ .
 For two different search radius. ο and 1's. this implies that From our observational results we infer that ATym(50.30)=L5 and AfZyw(100.50)=L8.," For two different search radius, $r_{\rm p,1}^{\rm max}$ and $r_{\rm p,2}^{\rm max}$ this implies that From our observational results we infer that $\Delta T_{\rm MM} (50,30) = 1.5$ and $\Delta T_{\rm MM} (100,50) = 1.8$."
 These values compare nicely with the ratios from Table 7 timescales. AT(50.30)=1.6 and ATym(100.50)2 1.8.," These values compare nicely with the ratios from Table \ref{tmm} timescales, $\Delta T_{\rm MM} (50,30) = 1.6$ and $\Delta T_{\rm MM} (100,50) = 1.8$ ."
 This supports the robustness of the assumed μμ. although the normalization of these timescales have a factor of two uncertainty.," This supports the robustness of the assumed $T_{\rm MM}$, although the normalization of these timescales have a factor of two uncertainty."
 We estimate the final major merger rate averaging the values derived from the 30. 50 and 1005! kpe merger fractions. and its error as the average of the individual merger rates” errors.," We estimate the final major merger rate averaging the values derived from the 30, 50 and $100h^{-1}$ kpc merger fractions, and its error as the average of the individual merger rates' errors."
 We obtain the minor merger rate. defined as the merger rate of 1/10€ui«1/4 close pairs. from the major one as where the factor ‘T accounts for the difference in the minor merger timescale with respect to the major merger one in close pars. Thin'YxἘν.," We obtain the minor merger rate, defined as the merger rate of $1/10 \leq \mu < 1/4$ close pairs, from the major one as where the factor $\Upsilon$ accounts for the difference in the minor merger timescale with respect to the major merger one in close pairs, $T_{\rm mm} = \Upsilon \times T_{\rm MM}$."
 Only a few studies in the literature attempt to estimate YT: study the merger timescale of dark matter haloes. finding f-2.," Only a few studies in the literature attempt to estimate $\Upsilon$: study the merger timescale of dark matter haloes, finding $\Upsilon \sim 2$."
 On the other hand. obtain 'Y=1.5+0.1 from N-body/hydrodynamical simulations.," On the other hand, obtain $\Upsilon = 1.5 \pm 0.1$ from $N$ -body/hydrodynamical simulations."
 As we have alreadyshown. the major merger timescales from are similar to ours. so we assume the minor-to-major merger time scale from in the following.," As we have alreadyshown, the major merger timescales from are similar to ours, so we assume the minor-to-major merger time scale from in the following."
" We also assume that the factor C,, for minor mergers is the same as the one for major mergers.", We also assume that the factor $C_{\rm m}$ for minor mergers is the same as the one for major mergers.
" Finally. the total merger rate is Ri,=Rum+Ry."," Finally, the total merger rate is $R_{\rm m} = R_{\rm MM} + R_{\rm mm}$."
 We summarize our results on the merger rates in Table 8.. and we show them in the Fig. 8..," We summarize our results on the merger rates in Table \ref{mrtab}, and we show them in the Fig. \ref{mrfig}."
 We find that We apply the steps in the previous section to estimate the major. minor and total merger rate of red and blue galaxies.," We find that We apply the steps in the previous section to estimate the major, minor and total merger rate of red and blue galaxies."
 We take Ti=3-9 Gyr and Tis=4.8 Gyr for r=ΙΟΟΗΓΙ kpe because of the different average stellar mass of red and blue principal galaxies. while the factor Y does not depend on the gas content ofthe galaxies(?).," We take $T_{\rm MM}^{\rm red} = 3.9$ Gyr and $T_{\rm MM}^{\rm blue} = 4.8$ Gyr for $r_{\rm p}^{\rm max} = 100h^{-1}$ kpc because of the different average stellar mass of red and blue principal galaxies, while the factor $\Upsilon$ does not depend on the gas content ofthe galaxies."
. The merger rates that we obtain are listed in Table 8.., The merger rates that we obtain are listed in Table \ref{mrtab}.
" The merger rates (minor and major) of red galaxies do not evolve with redshift in the range under study. RSS=0.064 Gyr!and Ri,=0.091 Gyr "," The merger rates (minor and major) of red galaxies do not evolve with redshift in the range under study, $R_{\rm mm}^{\rm red} = 0.064$ $^{-1}$and $R_{\rm MM}^{\rm red} = 0.091$ $^{-1}$ ."
find that the minor and major merger rate of Elliptical Like !.Objects (ELOs) atz~0.75 in their cosmological simulations areAj=0.06 Gyr! and Ry=0.08 Gyr7!. in good agreement with our observed values.," find that the minor and major merger rate of Elliptical Like Objects (ELOs) at $z \sim 0.75$ in their cosmological simulations are$R_{\rm mm} = 0.06$ $^{-1}$ and $R_{\rm MM} = 0.08$ $^{-1}$ , in good agreement with our observed values."
 On the other hand. model predicts that Run~Rum forfun=1/3 (seealso ?2)).whilefrom our observations we infer Ry=1.1xRuin for μμ= 1/3.," On the other hand, model predicts that $R_{\rm mm} \sim R_{\rm MM}$ for$\mu_{\rm MM} = 1/3$ (seealso ),whilefrom our observations we infer $R_{\rm mm} = 1.1 \times R_{\rm MM}$ for $\mu_{\rm MM} = 1/3$ ."
" The minor merger rate of blue galaxies. denoted Rh, increases by ~20% from z=0.8 to z2 0.5. butthe measured values are compatible with a constant merger rate within error"," The minor merger rate of blue galaxies, denoted $R_{\rm mm}^{\rm blue}$ , increases by $\sim$ from $z = 0.8$ to $z = 0.5$ , butthe measured values are compatible with a constant merger rate within error"
7? and ? have suggested that the apparent correlations found in mass-period ancl surface gravitv-period. plots for the known planetary systems may be due to evaporation processes.,\cite{mazeh} and \cite{pjw} have suggested that the apparent correlations found in mass-period and surface gravity-period plots for the known planetary systems may be due to evaporation processes.
 We have shown that if planets like the known sample existed. closer to their stars they would not have survived the saturated phase ofstellar X-ray/EUV emission. and hence would appear as à void in these plots.," We have shown that if planets like the known sample existed closer to their stars they would not have survived the saturated phase of stellar X-ray/EUV emission, and hence would appear as a void in these plots."
" If N-rav/IEUNV. irradiation is important. then the underlying cut-olf must be linear in the AL""o versus q-7 lane. as the ratio of potential energy to irraciating energv vcr unit time defines the planets characteristic lifetime."," If X-ray/EUV irradiation is important, then the underlying cut-off must be linear in the $M_p^2/R_p^3$ versus $a^{-2}$ plane, as the ratio of potential energy to irradiating energy per unit time defines the planet's characteristic lifetime."
 A xot of this formi is shown in Figure 6.., A plot of this form is shown in Figure \ref{printcorr}.
 A sparsely populated region is seen in the lower Left. similar to those in. Figure l.," A sparsely populated region is seen in the lower left, similar to those in Figure \ref{correlations}."
 For a given class of star it is possible to plot a power aw on this diagram. representing a destruction limit. below which planets would either have evaporated. completely. or eft planetary cores that are not. currently detectable. in ransit surveys.," For a given class of star it is possible to plot a power law on this diagram, representing a destruction limit, below which planets would either have evaporated completely or left planetary cores that are not currently detectable in transit surveys."
" The functional form of this destruction limit is where M, is the mass of the planet. Hy is the planetary radius. e is the semi major axis. L7"" is the saturated X-rav/EUVM luminosity of the star. Τμ is the saturation timescale of the spectral type and G is the gravitational constant."," The functional form of this destruction limit is where $_p$ is the mass of the planet, $_p$ is the planetary radius, $a$ is the semi major axis, $_{x}^{sat}$ is the saturated X-ray/EUV luminosity of the star, $\tau_{sat}$ is the saturation timescale of the spectral type and G is the gravitational constant."
 As can be seen in Figure 6.. the plotted lines corresponding to Εαν stars track the edge of the missing region very well without anv fitting.," As can be seen in Figure \ref{printcorr}, the plotted lines corresponding to FGK stars track the edge of the missing region very well without any fitting."
 All svstems exist above the destruction line for their spectral type., All systems exist above the destruction line for their spectral type.
 Uncertainties and assumptions in our simple evaporation model. (as discussed. in Sections 2.2.— and 2.3)) could move the destruction limits plotted in Fig. 6...," Uncertainties and assumptions in our simple evaporation model (as discussed in Sections \ref{evapmod} and \ref{xrayest}) ) could move the destruction limits plotted in Fig. \ref{printcorr}. .,"
 but they are unlikely to change the functional form. which matches the observed distribution well.," but they are unlikely to change the functional form, which matches the observed distribution well."
 The estruction limits can be raised bv factors. of SN. 3 and 8.5 for P. €G and [x stars. respectively before known planets (in their current state) fall into the danger zone.," The destruction limits can be raised by factors of 8, 3 and 8.5 for F, G and K stars respectively before known planets (in their current state) fall into the danger zone."
 ‘This suggests that our assumptions have tended: to underestimate the effects of evaporation., This suggests that our assumptions have tended to underestimate the effects of evaporation.
 This is consistent with our expectations. since neglecting Roche geometry. planet radius evolution (Sect. 2.2)).," This is consistent with our expectations, since neglecting Roche geometry, planet radius evolution (Sect. \ref{evapmod}) ),"
 X-ray irradiation since the saturated phase. ancl the observed excess N-ray emission of planet hosting stars (Sect. 2.3))," X-ray irradiation since the saturated phase, and the observed excess X-ray emission of planet hosting stars (Sect. \ref{xrayest}) )"
 are all likely to result in underestimates of evaporation rates., are all likely to result in underestimates of evaporation rates.
 We note that uncertainties in saturation timescales alone cannot move the destruction limits up to the known planets., We note that uncertainties in saturation timescales alone cannot move the destruction limits up to the known planets.
 This would require timescales of 0.8. 0.6 and 3 ινε [or E. €i and. Ix stars respectively. which are unreasonably high values compared. with existing X-ray observations (Sect. 2.3))," This would require timescales of 0.8, 0.6 and 3 Gyr for F, G and K stars respectively, which are unreasonably high values compared with existing X-ray observations (Sect. \ref{xrayest}) )."
can produce possible Bss with masses four times those of stars at the AIS turnoll (Leonard Linnell 1992).,can produce possible BSs with masses four times those of stars at the MS turnoff (Leonard Linnell 1992).
 Stellar collisions are also an important formation channel of BS formation in globular clusters (GC's) and open clusters (OCs: Cilebbeck et al., Stellar collisions are also an important formation channel of BS formation in globular clusters (GCs) and open clusters (OCs; Glebbeek et al.
 2008)., 2008).
 Given the high luminosities and common presence of BSs in stellar svstems (e.g. Ahumada Lapasset 1995 for OC's: Piotto et al.," Given the high luminosities and common presence of BSs in stellar systems (e.g., Ahumada Lapasset 1995 for OCs; Piotto et al."
 2002 for GCs: Mapelli et al., 2002 for GCs; Mapelli et al.
 2009 for clwarl galaxies). we believe that we must consider the elfects of BSs in studies of stellar populations using population synthesis. applied to unresolved: observations.," 2009 for dwarf galaxies), we believe that we must consider the effects of BSs in studies of stellar populations using population synthesis applied to unresolved observations."
 The key issue is how to accurately include BS contributions in SSP models., The key issue is how to accurately include BS contributions in SSP models.
 Studies show that no single mechanism can account or the entire DS. population observed in any one star cluster (Stryker 1993)., Studies show that no single mechanism can account for the entire BS population observed in any one star cluster (Stryker 1993).
 This means that it is not casy o theoretically measure the respective contribution of Biss in SSPs of dillerent ages and metallicities., This means that it is not easy to theoretically measure the respective contribution of BSs in SSPs of different ages and metallicities.
 Pherclore. xulding up BS population characteristics empirically [ron he statistics of a large sample of star clusters. could. be more practical and reliable than relying on. incomplete heoretical approaches.," Therefore, building up BS population characteristics empirically from the statistics of a large sample of star clusters could be more practical and reliable than relying on incomplete theoretical approaches."
 This way. the behaviour of BSs (in erms of their specific frequency and. relative distribution with respect to the ALS turnoll in CMDs) can be mocdelled.," This way, the behaviour of BSs (in terms of their specific frequency and relative distribution with respect to the MS turnoff in CMDs) can be modelled."
 OCs in the Galaxy have à number of advantages for use as a working sample. for instance. (i) they are good observational templates of ideal SSPs in the real world. and all DSs in a given OC belong to a single population: (ii) many OC's have mult-epoch proper-motion and/or radial-velocitv data. so that their cluster membership probabilities can be measured accurately.," OCs in the Galaxy have a number of advantages for use as a working sample, for instance, (i) they are good observational templates of ideal SSPs in the real world, and all BSs in a given OC belong to a single population; (ii) many OCs have multi-epoch proper-motion and/or radial-velocity data, so that their cluster membership probabilities can be measured accurately."
 Preliminary modelling of BS effects has been done on the basis of individual clusters in our previous papers (Deng et al., Preliminary modelling of BS effects has been done on the basis of individual clusters in our previous papers (Deng et al.
 1999: Nin Deng 2005: Nin et al, 1999; Xin Deng 2005; Xin et al.
 2007. 2008). where we (i) introduced. the method. used. for calculating the realistic. integrated spectrum of a star. cluster. including the contribution of DBSs: (i) analysed the modifications to the integrated. spectra and broad-band colours caused by Bss: and (ii) estimated: the possible uncertainties in the conventional SSP models. showing that the ages of star clusters can be underestimated by up to if BS contributions are not considered.," 2007, 2008), where we (i) introduced the method used for calculating the realistic, integrated spectrum of a star cluster including the contribution of BSs; (ii) analysed the modifications to the integrated spectra and broad-band colours caused by BSs; and (iii) estimated the possible uncertainties in the conventional SSP models, showing that the ages of star clusters can be underestimated by up to if BS contributions are not considered."
 In this paper. we present a set of BS-SSP models based on a statistical study of Galactic OC's.," In this paper, we present a set of BS-SSP models based on a statistical study of Galactic OCs."
 A working sample including LOO Galactic OC's is used. to reduce stochastic ellects in collecting the required. statistical relations., A working sample including 100 Galactic OCs is used to reduce stochastic effects in collecting the required statistical relations.
 We also use the standard. deviation (0) to keep the statistics in a relatively small dispersion region., We also use the standard deviation $\sigma$ ) to keep the statistics in a relatively small dispersion region.
 Our mocoels cover the wavelength. range from 91 to 160. jun. ages from 0.1 to 20 Gaver and metallicities Z=0.0004.0.004.0.008.0.02 (solar metallicity) ane 0.05.," Our models cover the wavelength range from 91 to 160 $\mu$ m, ages from 0.1 to 20 Gyr and metallicities $Z=0.0004, 0.004,
0.008, 0.02$ (solar metallicity) and 0.05."
 This paper is organised. as follows., This paper is organised as follows.
 In Section 2. we present the statistical results of the properties. of DS populations based on 100 Calactie OC's.," In Section 2, we present the statistical results of the properties of BS populations based on 100 Galactic OCs."
 Here. we also describe the construction procedure of our. models.," Here, we also describe the construction procedure of our models."
 In Section 3. we discuss our mocel results ancl compare them with those of DCO3.," In Section 3, we discuss our model results and compare them with those of BC03."
 In Section d. we test our models bv comparison with the broad-band. colours of OC's.," In Section 4, we test our models by comparison with the broad-band colours of OCs."
 Our model colours are in good agreement: with observed. OC colours (which include BS contributions)., Our model colours are in good agreement with observed OC colours (which include BS contributions).
 In Section 5. we test our. models using observed. integrated. spectra of star clusters.," In Section 5, we test our models using observed integrated spectra of star clusters."
 Compared to BCOS. our models vield more accurate age predictions at a range of wavelengths. including in the optical regime.," Compared to BC03, our models yield more accurate age predictions at a range of wavelengths, including in the optical regime."
 Finally. a summary. and. further brief discussion is contained in Section 6.," Finally, a summary and further brief discussion is contained in Section 6."
 The fundamental ingredients. of our models are [listed in Table 1.., The fundamental ingredients of our models are listed in Table \ref{table1}.
 For convenience. we use the widely adopted ὃς09. models as reference.," For convenience, we use the widely adopted BC03 models as reference."
 Mocifications owing to BSs are calculated. as increments to the DCOS SSPs of the same age and metallicity., Modifications owing to BSs are calculated as increments to the BC03 SSPs of the same age and metallicity.
 We adopted the Padoval094 isochrones (Bertelli ct al., We adopted the Padova1994 isochrones (Bertelli et al.
 1994) and. the BasSeL spectral library (Lejeune et al., 1994) and the BaSeL spectral library (Lejeune et al.
 L997) because they homogencously cover the widest ranges of age and metallicity. anc the longes wavelength range.," 1997) because they homogeneously cover the widest ranges of age and metallicity, and the longest wavelength range."
 The Pacdova2000 isochrones (Cirardi ο al., The Padova2000 isochrones (Girardi et al.
 2000) are based on a more recent equation of state ancl low-temperature opacities compared. to. Padoval094., 2000) are based on a more recent equation of state and low-temperature opacities compared to Padova1994.
 Llowever. we decided against adopting them for our mode construction. because BCO3 do not recommend to use their SSPs based on the Padova2000 isochrones.," However, we decided against adopting them for our model construction, because BC03 do not recommend to use their SSPs based on the Padova2000 isochrones."
" They state tha their models based. on the Pacdova2000 isochrones ""tene to produce worse agreement with observed galaxy colours? BeOS. their. footnote. 6)."," They state that their models based on the Padova2000 isochrones “tend to produce worse agreement with observed galaxy colours” (BC03, their footnote 6)."
 High-resolution. observationa spectral libraries (e.g... Pickles et al.," High-resolution observational spectral libraries (e.g., Pickles et al."
 1998: Le Borgne e al., 1998; Le Borgne et al.
 2003) suller from problems related to limited: parameter coverage., 2003) suffer from problems related to limited parameter coverage.
 Instead. of combining spectra. from dilferen libraries to enlarge our parameter coverage. we decided to use onlv the theoretical library.," Instead of combining spectra from different libraries to enlarge our parameter coverage, we decided to use only the theoretical library."
 For consistency. with BCOS. we adopted the Salpeter and. Chabrier stellar initial mass functions (AALS: Salpeter 1955: Chabrier 2003).," For consistency with BC03, we adopted the Salpeter and Chabrier stellar initial mass functions (IMFs; Salpeter 1955; Chabrier 2003)."
 As clearly shown by Dabringhausen et al. (, As clearly shown by Dabringhausen et al. (
2008. their fig.,"2008, their fig."
 S). the Chabrier ΙΔΗΣ is. almost inclistinguishable from the Kroupa EME (Ixroupa 2001. 2002) when normalised as Fronn=JM. which is exactly how both BCO3 and we ourselves normalise the SSP models.," 8), the Chabrier IMF is almost indistinguishable from the Kroupa IMF (Kroupa 2001, 2002) when normalised as $\int_{0.1}^{100}\xi(m){\rm
 d}m=1{\rm M}_{\odot}$, which is exactly how both BC03 and we ourselves normalise the SSP models."
 Using such a normalisation. the slight dillerences between the two LIMES cannot cause any ellective modifications as regards the BS contribution to SSPs.," Using such a normalisation, the slight differences between the two IMFs cannot cause any effective modifications as regards the BS contribution to SSPs."
 Therefore. we refer to the IME (in addition to the Salpeter ΠΟ) as the ‘Canonical," Therefore, we refer to the IMF (in addition to the Salpeter IMF) as the `Canonical"
[εαν ane fous in order to get the dispersion relation.,$f_{xxy}$ and $f_{xxz}$ in order to get the dispersion relation.
 These can be derived to be: where we neglect terms that are of higher order than f;. being its ὃν component and fj.," These can be derived to be: where we neglect terms that are of higher order than $\fv_3$ , being its $\partial_t$ component and $\fv_4$."
 I5qs., Eqs.
 29. are used to solve for fey and fos. which can be substituted in the equation for fi to give the mioclified dispersion relation: Fig.," \ref{eq:dtf3} are used to solve for $f_{xy}$ and $f_{xz}$, which can be substituted in the equation for $\fv_1$ to give the modified dispersion relation: Fig."
"e 6 shows the ogrowth rate for the dispersion1 relation up to £f; as a function of both v/a, and &.", \ref{fig:wf1f3} shows the growth rate for the dispersion relation up to $\fv_3$ as a function of both $\nu/\omega_g$ and $k$.
 On the long scales there is no noticeable dillerence between this result and the one ignoring f;., On the long scales there is no noticeable difference between this result and the one ignoring $\fv_3$.
 Phe peak around resonance softens a bit and moves to slightly shorter wave numbers. closer to Airy=1.," The peak around resonance softens a bit and moves to slightly shorter wave numbers, closer to $kr_g=1$."
 We have evaluated the dominant terms for subsequent orders in f (up to fj5) and found that this trend seems to continue when higher orders of f are included., We have evaluated the dominant terms for subsequent orders in $\fv$ (up to $\fv_{10}$ ) and found that this trend seems to continue when higher orders of $\fv$ are included.
 H£ the trend continues as we expect. the transition around. resonance smoothens out completely.," If the trend continues as we expect, the transition around resonance smoothens out completely."
 A curious οσοι arises for the left-hand mode in. the short-wavelength limit., A curious effect arises for the left-hand mode in the short-wavelength limit.
 Rather than steeply declining with wave number. it increases when the highest order of f that is included. is odd. (ic. f; or £2).," Rather than steeply declining with wave number, it increases when the highest order of $\fv$ that is included is odd (i.e. $\fv_3$ or $\fv_5$ )."
 When the highest order of f that is included is an even number. the growth rate asymptotically decreases with wave number.," When the highest order of $\fv$ that is included is an even number, the growth rate asymptotically decreases with wave number."
 From the dispersion relation it can be seen how this ellect arises: When the highest order is an even number. the orders of & in the nominator are exactly matehect in the denominator and thus cancel. whereas this is not the case when the highest order is an odd number.," From the dispersion relation it can be seen how this effect arises: When the highest order is an even number, the orders of $k$ in the nominator are exactly matched in the denominator and thus cancel, whereas this is not the case when the highest order is an odd number."
 In the analvsis by ? the growth rate in this regime is shown to decline for high &. which leacs us to believe that the real trend is captured when the highest order is an even order of f (fe. £j). which has a decreasing erowth rate with wave number for the left-hand mocdoe.," In the analysis by \citet{1983Achterberg} the growth rate in this regime is shown to decline for high $k$, which leads us to believe that the real trend is captured when the highest order is an even order of $\fv$ $\fv_2$, $\fv_4$ ), which has a decreasing growth rate with wave number for the left-hand mode."
 We have derived the dispersion relation for an instability hat is driven by the zeroth order cosmic ray current., We have derived the dispersion relation for an instability that is driven by the zeroth order cosmic ray current.
 The ong-wavelength component is mediated by the stress tensor. and coupling between the left- and right-hand modes occurs when scattering of cosmic ravs olf small scale turbulence is aken into account.," The long-wavelength component is mediated by the stress tensor, and coupling between the left- and right-hand modes occurs when scattering of cosmic rays off small scale turbulence is taken into account."
 ‘The streaming of cosmic ravs ahead of the shock induces a return current in the plasma that triggers instabilities from arge scales down to the Alfvénnic regime on very small scales at which the magnetic tension becomes ellicient. in damping the oscillations., The streaming of cosmic rays ahead of the shock induces a return current in the plasma that triggers instabilities from large scales down to the Alfvénnic regime on very small scales at which the magnetic tension becomes efficient in damping the oscillations.
 We did not include the Alfvénn erm (magnetic tension). which only becomes of importance ond Aryc107 for our parameters.," We did not include the Alfvénn term (magnetic tension), which only becomes of importance beyond $k r_g > 10^3$ for our parameters."
 On small. scales our analvsis reproduces the hiverodyvnamical non-resonant instability as described in ?.., On small scales our analysis reproduces the hydrodynamical non-resonant instability as described in \citet{2004Bell}.
 The growth rate at small scales is very high anc the amplified. magnetic field. can reach many times its zeroth order strength., The growth rate at small scales is very high and the amplified magnetic field can reach many times its zeroth order strength.
 In an environment such as that upstream of a SNR. blast wave. the ellicient. gencration of magnetic field amplification on small scales is expected to increase the scattering ellicienev of the lower energy cosmic ravs.," In an environment such as that upstream of a SNR blast wave, the efficient generation of magnetic field amplification on small scales is expected to increase the scattering efficiency of the lower energy cosmic rays."
 We find that this in turn. seems to couple the Ieft- and right-hand. modes. resulting in a small left-hancecdlvy polarised contribution to. the short-scale instability. in addition to the dominant right-handed moce.," We find that this in turn seems to couple the left- and right-hand modes, resulting in a small left-handedly polarised contribution to the short-scale instability, in addition to the dominant right-handed mode."
 On long scales. predominantlv the left-hand mode is unstable.," On long scales, predominantly the left-hand mode is unstable."
 This effect. only arises when anisotropy of order f» (and higher) istaken into account., This effect only arises when anisotropy of order $\fv_2$ (and higher) istaken into account.
 The long- mode behaviour already. converges when orders of Ἐν are included. as we find that higher-order anisotropy terms do not alter the instability. significantly.," The long-wavelength mode behaviour already converges when orders of $\fv_2$ are included, as we find that higher-order anisotropy terms do not alter the instability significantly."
 Inclusion, Inclusion
the case of GX 354-0 where a~ 1.6.,the case of GX 354–0 where $\alpha \sim$ 1.6.
 We give a possible interpretation of these results in the Discussion., We give a possible interpretation of these results in the Discussion.
 The determination of spectral index « obtained from Comptonization of seed photons in a bounded medium has been faced since a long time., The determination of spectral index $\alpha$ obtained from Comptonization of seed photons in a bounded medium has been faced since a long time.
 The emerging radiation spectrum depends on several parameters such as the geometry of the plasma (e.g.. slab or sphere). the electron temperature and optical depth. and the space distribution. of the seed photons inside the medium.," The emerging radiation spectrum depends on several parameters such as the geometry of the plasma (e.g., slab or sphere), the electron temperature and optical depth, and the space distribution of the seed photons inside the medium."
 Sunyaev Titarchuk (1980) report the value o obtained from the solution of the stationary radiative transfer equation in the non-relativistic case (Fokker-Planxk approximation) obtained as a convolution of the time-dependent equation with the time-escape probability distribution P(u) for the case of source photons distributed according to the first eigenfunction of the space operator., Sunyaev Titarchuk (1980) report the value $\alpha$ obtained from the solution of the stationary radiative transfer equation in the non-relativistic case (Fokker-Planxk approximation) obtained as a convolution of the time-dependent equation with the time-escape probability distribution P(u) for the case of source photons distributed according to the first eigenfunction of the space operator.
 Later Titarchuk.(1994).. Hua&Titarchuk(1995) and TL95 extended the results to the sub-relativistic case considering both slab and spherical geometry.," Later \cite{t94}, \cite{ht95} and TL95 extended the results to the sub-relativistic case considering both slab and spherical geometry."
 In order to understand what it does happen in NS LMXBs sources. one has to consider the hydrodynamical conditions in the region between the accretion disk and the NS surface.," In order to understand what it does happen in NS LMXBs sources, one has to consider the hydrodynamical conditions in the region between the accretion disk and the NS surface."
 We will refer to this region as the transition layer (TL). often also called boundary layer.," We will refer to this region as the transition layer (TL), often also called boundary layer."
 Actually. the production of a strong TC bump in the persistent X-ray spectra of NS LMXBs is thought to originate in this TL. namely the region where matter deviates from its Keplerian angular velocity in order to match that of the slowly spinning NS.," Actually, the production of a strong TC bump in the persistent X-ray spectra of NS LMXBs is thought to originate in this TL, namely the region where matter deviates from its Keplerian angular velocity in order to match that of the slowly spinning NS."
 The radiative and hydrodynamical configuration of the TL is mostly dictated by the Reynolds number *1999).. which is proportional to the mass accretion rate and is eventually the inverse of the viscosity parameter of the Shakura-Sunyaev disk.," The radiative and hydrodynamical configuration of the TL is mostly dictated by the Reynolds number $\gamma$, which is proportional to the mass accretion rate and is eventually the inverse of the viscosity parameter of the Shakura-Sunyaev disk."
 In particular. + ddetermines the radial extension of the TL and in turn. from the mass accretio rate. its optical depth.," In particular, $\gamma$ determines the radial extension of the TL and in turn, from the mass accretion rate, its optical depth."
 It is worth pointing out that determination of the vertical height-scale of the TL is in fact à very complicated problem. as it should require a complete 3D magneto-hydrodynamical treatment of the problem.," It is worth pointing out that determination of the vertical height-scale of the TL is in fact a very complicated problem, as it should require a complete 3D magneto-hydrodynamical treatment of the problem."
 Using the slim disk (thus vertical-averaged) equations for determining the radial thermo-hydrodynamical structure of the TL may be in fact an issue (e.g..Popham&Sunyaev2001).," Using the slim disk (thus vertical-averaged) equations for determining the radial thermo-hydrodynamical structure of the TL may be in fact an issue \citep[e.g.,][]{ps01}."
". The enhanced radiation. and thermal pressure because of higher electron temperature are expected to increase the vertical height-scale of the TL with Η ~&,,,.", The enhanced radiation and thermal pressure because of higher electron temperature are expected to increase the vertical height-scale of the TL with H $\approx R_{ns}$.
 Moreover. the solution for the angular momentum equation (Titarchuketal.1998:Titarchuk&Osherovich1999). for Reynolds number 5510 gives a TL radial extension Aer;0.5Rys.," Moreover, the solution for the angular momentum equation \citep{tlm98,to99} for Reynolds number $\gamma > 5-10$ gives a TL radial extension $\Delta R_{TL} \la 0.5 R_{NS}$."
 With these characteristic length-scales. it seems more plausible to approximate the TL geometry to à slab whose normal is directed along the disk plane.," With these characteristic length-scales, it seems more plausible to approximate the TL geometry to a slab whose normal is directed along the disk plane."
 It is worth emphasizing that Haardt Maraschi (1993. hereafter HM93) determined the theoretical Comptonization index à considering a two-phase model for accretion disks in AGN in which a hot corona ts surrounding and sandwiching the underlying cold accretion disk.," It is worth emphasizing that Haardt Maraschi (1993, hereafter HM93) determined the theoretical Comptonization index $\alpha$ considering a two-phase model for accretion disks in AGN in which a hot corona is surrounding and sandwiching the underlying cold accretion disk."
 The model can in principle be applied also to the case of solar-mass BH sources., The model can in principle be applied also to the case of solar-mass BH sources.
 HM93 assume that the corona and the disk are two slabs at significant different temperatures and put in contact each other., HM93 assume that the corona and the disk are two slabs at significant different temperatures and put in contact each other.
 The authors concentrated on the case of high temperature (AZ).& 50 keV) and low optical depth (7<1) for the corona. so that the diffusion approximation cannot hold. unlike what we are considering.," The authors concentrated on the case of high temperature $\kte \ga $ 50 keV) and low optical depth $\tau < $ 1) for the corona, so that the diffusion approximation cannot hold, unlike what we are considering."
 One of the consequences of the high-temperature treatment is that electron scattering is anisotropic. with a significant fraction of the power back- the disk., One of the consequences of the high-temperature treatment is that electron scattering is anisotropic with a significant fraction of the power back-irradiating the disk.
" In the HM93 model. the boundary condition of the hot corona is the disk cool surface (with ATi, <5 eV) with energy-dependent albedo."," In the HM93 model, the boundary condition of the hot corona is the disk cool surface (with $\ktbb < $ 5 eV) with energy-dependent albedo."
 Note also that in their geometry. of the disk flux is intercepted and reprocessed by the top plasma.," Note also that in their geometry, of the disk flux is intercepted and reprocessed by the top plasma."
 In the geometry considered here on the other hand. it is possible that part of the disk emission directly escapes the system. while a fraction of its flux is intercepted by the TL.," In the geometry considered here on the other hand, it is possible that part of the disk emission directly escapes the system, while a fraction of its flux is intercepted by the TL."
 We are actually interested here on this portion of the intercepted disk flux (see next Section)., We are actually interested here on this portion of the intercepted disk flux (see next Section).
 Because of these differences between the two models. a direct comparison of the derived theoretical results 1s not straightforward.," Because of these differences between the two models, a direct comparison of the derived theoretical results is not straightforward."
 The reader can refer to the paper of HM93 for further details in order to better understand the differences of our and their approaches., The reader can refer to the paper of HM93 for further details in order to better understand the differences of our and their approaches.
 The energy balance in the TL ts dictated by Coulomb collisions with protons (gravitational energy release). while inverse Compton and free-free emission are the main cooling channels (see a formulation of this problem in the pioneer work by Zel'dovich Shakura 1969).," The energy balance in the TL is dictated by Coulomb collisions with protons (gravitational energy release), while inverse Compton and free-free emission are the main cooling channels (see a formulation of this problem in the pioneer work by Zel'dovich Shakura 1969)."
 In fact. for the characteristics electron temperature (3 keV <AL.x; 30 keV) and density," In fact, for the characteristics electron temperature (3 keV $\la kT_e\la$ 30 keV) and density"
"to be the sum of the angular momenta of all disc particles that are deemed part of the final disc remnant, calculated in the reference frame of the binary.","to be the sum of the angular momenta of all disc particles that are deemed part of the final disc remnant, calculated in the reference frame of the binary."
" While each set has some higher-density regions in this plane, in general most combinations of these angles are represented."," While each set has some higher-density regions in this plane, in general most combinations of these angles are represented."
" We have highlighted all runs that have a mostly orthogonal orientation between the disc angular momentum and v;«, which we take as within 20 degrees of orthogonal, with boxes."," We have highlighted all runs that have a mostly orthogonal orientation between the disc angular momentum and $v_{rel}$, which we take as within 20 degrees of orthogonal, with boxes."
 The blue box encompasses all nearly-orthogonal runs that also have the disc and the binary aligned to within 30 degrees; the red box encompasses all other disc-binary orientations., The blue box encompasses all nearly-orthogonal runs that also have the disc and the binary aligned to within 30 degrees; the red box encompasses all other disc–binary orientations.
" In the smaller panels of the figure, we trace back to the initial conditions to see if these interesting boxes are populated from a specific region of parameter space, plotting the density of runs that end up populating each box."," In the smaller panels of the figure, we trace back to the initial conditions to see if these interesting boxes are populated from a specific region of parameter space, plotting the density of runs that end up populating each box."
" The only interesting features turns out to be in the r,— 0 plane.", The only interesting features turns out to be in the $r_p$ --cos $\theta$ plane.
" The first notable result is that the (10,3)+10 series contains no runs that led to alignment between the binary and a disc that is orthogonal to the relative velocity between the single and the binary."," The first notable result is that the (10,3)+10 series contains no runs that led to alignment between the binary and a disc that is orthogonal to the relative velocity between the single and the binary."
" The equal mass series, in contrast, has several runs in this category."," The equal mass series, in contrast, has several runs in this category."
" Crucially, the initial conditions these runs came from are much smaller than the total parameter range we covered."," Crucially, the initial conditions these runs came from are much smaller than the total parameter range we covered."
 Roughly 10 per cent of all initial conditions with periastra less than about 20 au and prograde encounter angles result in high stellar velocities and remnant discs with similar orientations to Source I. All of our simulations have taken the binary as initially circular., Roughly 10 per cent of all initial conditions with periastra less than about 20 au and prograde encounter angles result in high stellar velocities and remnant discs with similar orientations to Source I. All of our simulations have taken the binary as initially circular.
 We can draw on extensive binary-single scattering work to estimate the effect of this simplifying assumption., We can draw on extensive binary–single scattering work to estimate the effect of this simplifying assumption.
" presents scattering results for equal-mass systems at three eccentricities: e—0, 0.7, and 0.99."," presents scattering results for equal-mass systems at three eccentricities: $e=0$, 0.7, and 0.99."
" He shows that the cross sections for scattering results (see the appendix for some discussion of this) depend only weakly on the eccentricity; increasing the eccentricity from 0 to 0.7 approximately doubles the cross section for both simple exchange encounters and resonant encounters, and a further increase to 0.99 has minimal effect."," He shows that the cross sections for scattering results (see the appendix for some discussion of this) depend only weakly on the eccentricity; increasing the eccentricity from 0 to 0.7 approximately doubles the cross section for both simple exchange encounters and resonant encounters, and a further increase to 0.99 has minimal effect."
" This increase in encounter cross section is mainly manifested in encounters that only slightly change the binary energy, however1984)."," This increase in encounter cross section is mainly manifested in encounters that only slightly change the binary energy, however."
". Over the range of velocities we consider, the cross section for a given energy change to the binary as a result of a scattering encounter in the range of interest, i.e. those encounters that generate Όκει>20s!, are almost identical for all three eccentricities."," Over the range of velocities we consider, the cross section for a given energy change to the binary as a result of a scattering encounter in the range of interest, i.e. those encounters that generate $v_{rel}>20$, are almost identical for all three eccentricities."
" T'he total number of sufficiently velocity-boosting encounters is thus largely insensitive to the eccentricity of the binary, and our zero-eccentricity assumption should not affect these results."," The total number of sufficiently velocity-boosting encounters is thus largely insensitive to the eccentricity of the binary, and our zero-eccentricity assumption should not affect these results."
" We also assume that the binary and its disc are largely dynamically unprocessed, in the sense that the disc is coplanar to the binary."," We also assume that the binary and its disc are largely dynamically unprocessed, in the sense that the disc is coplanar to the binary."
" Our results do rely on this assumption; if the disc and the binary are seriously misaligned prior to the encounter, our effort to trace back the interesting outcomes to their initial conditions (figure 9)) is compromised."," Our results do rely on this assumption; if the disc and the binary are seriously misaligned prior to the encounter, our effort to trace back the interesting outcomes to their initial conditions (figure \ref{angles}) ) is compromised."
" While there would almost certainly be a set of encounters that are more likely to lead to the proper disc-v.e; orientation, it might not be as nicely contained as in the coplanar case."," While there would almost certainly be a set of encounters that are more likely to lead to the proper $v_{rel}$ orientation, it might not be as nicely contained as in the coplanar case."
" The initial parameters of our disc, for which we have attempted to choose plausible values, are also an assumption."," The initial parameters of our disc, for which we have attempted to choose plausible values, are also an assumption."
" If the original disc is considerably larger or more compact than our choice of 100 au, the fractional amount of disc material retained will be different than the values quoted here."," If the original disc is considerably larger or more compact than our choice of 100 au, the fractional amount of disc material retained will be different than the values quoted here."
" Because the bulk of the retained material is originally in the inner regions of the disc, the conclusion that non-negligible amounts of material can remain bound should hold true as long as the inner regions of the disc are still intact."," Because the bulk of the retained material is originally in the inner regions of the disc, the conclusion that non-negligible amounts of material can remain bound should hold true as long as the inner regions of the disc are still intact."
 The effect of tidal or direct interactions with low mass cluster members prior to the encounter with the BN, The effect of tidal or direct interactions with low mass cluster members prior to the encounter with the BN
From the curves of Fig. 3..,"From the curves of Fig. \ref{fig:fin_dist},"
 we computed a. the error in the vignetting factor when we apply Eq. (9))," we computed $\frac{\Delta V}{V}$, the error in the vignetting factor when we apply Eq. \ref{eq:vign}) )"
 to a Wolter-I profile. as in the definition of Eq. (8)).," to a Wolter-I profile, as in the definition of Eq. \ref{eq:max_diam}) )."
 This quantity is plotted in Fig., This quantity is plotted in Fig.
 4 for the case L’=0.5., \ref{fig:error_eval} for the case $L' = 0.5$.
 It ean be seen that. in the cases interesting for us. Le.. for sufficiently large f# and not too large values of 0. is positive and of the order of a few percent.," It can be seen that, in the cases interesting for us, i.e., for sufficiently large $f\#$ and not too large values of $\delta$, $\frac{\Delta V}{V}$ is positive and of the order of a few percent."
 We repeated this atexercise for several values of L’ in the interval 0.25-| to investigate the dependence of ar on this parameter. and it turned out that this ratio. if positive and if V(Wolter)«I. can be approximated very well by the empirical formula (see Fig.," We repeated this exercise for several values of $L'$ in the interval $0.25 - 1$ to investigate the dependence of $\frac{\Delta V}{V}$ on this parameter, and it turned out that this ratio, if positive and if $V(\mbox{Wolter}) < 1$, can be approximated very well by the empirical formula (see Fig."
 and Table 1)) where y«14.3 is with very good approximation a constant in the explored range of L’ values., \ref{fig:error_eval} and Table \ref{tab:DC_toler}) ) where $\gamma \approx 14.3$ is with very good approximation a constant in the explored range of $L'$ values.
 This equation has the same kind of dependence as Eq. (6)), This equation has the same kind of dependence as Eq. \ref{eq:angle_var1}) )
 for the slope variation alongthe profile of the mirror and Eq. (8)), for the slope variation alongthe profile of the mirror and Eq. \ref{eq:max_diam}) )
 for the area of the primary segment., for the area of the primary segment.
 For small £3 values. Eq. (12))," For small $f\#$ values, Eq. \ref{eq:F_var}) )"
 is not obeyed because of the saturation of V(Wolter) to I. so the expression of ar should be interpreted as an upper limit.," is not obeyed because of the saturation of $V$ (Wolter) to 1, so the expression of $\frac{\Delta V}{V}$ should be interpreted as an upper limit."
 Finally. for very large 6 the error deviates from Eq. (12))," Finally, for very large $\delta$ the error deviates from Eq. \ref{eq:F_var}) )"
 as it becomes negative. but this occurs only when V-0. so its weight in determining the effective area Is expected to be negligible.," as it becomes negative, but this occurs only when $V \rightarrow 0$, so its weight in determining the effective area is expected to be negligible."
 We can therefore conclude that the error introduced by the double cone approximation. regarding the effective area. is definitively smaller than L’/f.," We can therefore conclude that the error introduced by the double cone approximation, regarding the effective area, is definitively smaller than $L'/f\#$."
 In other words. the double cone approximation is valid when L-«f. a condition fulfilled in almost all practical cases.," In other words, the double cone approximation is valid when $L \ll f$, a condition fulfilled in almost all practical cases."
 Within the limits of this approximation. we derive in the next section the effective area for a source on- and off-axis.," Within the limits of this approximation, we derive in the next section the effective area for a source on- and off-axis."
 We assume that we can approximate the Wolter-I profile with a double cone profile. by adopting the tolerances estimated in Sect. 2..," We assume that we can approximate the Wolter-I profile with a double cone profile, by adopting the tolerances estimated in Sect. \ref{DC_WI}."
 [n the following. we adopt the convention to denote with Αρ.ϐ) the area of the mirror at the photon wavelength Jt. for a source off-axis by &. at a distance D (finite or infinite).," In the following, we adopt the convention to denote with $A_D(\lambda, \theta)$ the area of the mirror at the photon wavelength $\lambda$, for a source off-axis by $\theta$, at a distance $D$ (finite or infinite)."
 When referring to the area. we omit 2 and use the notation Ap(@).," When referring to the area, we omit $\lambda$ and use the notation $A_D(\theta)$."
 Firstly. we assume L;=£L» L.," Firstly, we assume $L_1 = L_2 = L$ ."
" The geometrical. collecting area of the primary segment. às seel by a source on-axis at ""infinite"" (1e... astronomical) distance is where we use the approximation Ay—Ro=~ Log."," The geometrical, collecting area of the primary segment, as seen by a source on-axis at “infinite"" (i.e., astronomical) distance is where we use the approximation $R_{\mathrm M} - R_0 \simeq L \alpha_0$ ."
 In this case. all reflected rays also undergo an identical reflection on the secondary segment (Fig. 1)).," In this case, all reflected rays also undergo an identical reflection on the secondary segment (Fig. \ref{fig:mirror_section}) ),"
 therefore Αι(0)=(0).," therefore $A_{\infty}(0)=A_{1,\infty}(0)$."
 We now denote by ία) the reflectivity of the mirror at the photon wavelength Jt. for a generic incidence angle a.," We now denote by $r_{\lambda}(\alpha)$ the reflectivity of the mirror at the photon wavelength $\lambda$, for a generic incidence angle $\alpha$."
 The form of this function depends on the coating structure: for a single layer coating. which operates in total external reflection. it slowly decreases up to the critical angle for Jt. followed by a sudden cutoff.," The form of this function depends on the coating structure: for a single layer coating, which operates in total external reflection, it slowly decreases up to the critical angle for $\lambda$, followed by a sudden cutoff."
 If a multilayer coating is used. ία) is a more complicated function and can be computed using one of the standard methods (e.g.. Parrat 1954:: Abeléss 1950)). by also including the etfect of roughness using. e.g.. the Croce (1980)) approach.," If a multilayer coating is used, $r_{\lambda}(\alpha)$ is a more complicated function and can be computed using one of the standard methods (e.g., Parrat \cite{Parrat}; Abelèss \cite{Abeles}) ), by also including the effect of roughness using, e.g., the N\'evvot-Croce \cite{NevotCroce}) ) approach."
 Since photons are reflected twice at the same angle αρ. we multiply the geometrical area by the squared reflectivity to obtain the area at the photon wavelength /1 This is à well known result.," Since photons are reflected twice at the same angle $\alpha_0$, we multiply the geometrical area by the squared reflectivity to obtain the area at the photon wavelength $\lambda$ This is a well known result."
 We now keep the source on- but at a finite distance D and assume more generally that Lj4Ls (with L=Ls as à particular case).," We now keep the source on-axis, but at a finite distance $D$ and assume more generally that $L_1 \neq L_2$ (with $L_1 = L_2$ as a particular case)."
 As already discussedin Sect. 2..," As already discussedin Sect. \ref{DC_WI},"
 all mirror sectors see the source off-axis by the same angle 0=Ro/D., all mirror sectors see the source off-axis by the same angle $\delta = R_0/D$.
" The effective area of the segment thereby becomes where «a,=«o+0.", The effective area of the segment thereby becomes where $\alpha_1 = \alpha_0+\delta$.
 The effective area for a source at finite distance is then obtained from Eq. (15)).," The effective area for a source at finite distance is then obtained from Eq. \ref{eq:Apar_fin_on}) ),"
 times the vignetting factor of Eq. (9)).," times the vignetting factor of Eq. \ref{eq:vign}) ),"
 times r)(a2). the reflectivity of the secondary mirror segment. where a2=ao—0./£ V<1.," times $r_{\lambda}(\alpha_2)$, the reflectivity of the secondary mirror segment, where $\alpha_2 = \alpha_0-\delta$. $V<1$,"
 substitution of Eq. (9)), substitution of Eq. \ref{eq:vign}) )
 yields As one might expect. the geometric area is that of the secondary segment. projected onto the wavefront after the primary reflection.," yields As one might expect, the geometric area is that of the secondary segment, projected onto the wavefront after the primary reflection."
 We note that by setting 0= we retrieve the on-axis result. in Eq. (14)).," We note that by setting $\delta = 0$ we retrieve the on-axis result, in Eq. \ref{eq:Ae_inf_on}) )."
 On the other hand. tf Lie.>Lyay. all the primary segment ts effective in the double reflection. κο V=1.," On the other hand, if $L_2\alpha_2 >L_1\alpha_1$, all the primary segment is effective in the double reflection, so $V =1$."
 This may occur with a divergent source on-axis. if lL.> Ld," This may occur with a divergent source on-axis, if $L_2 \gg L_1$."
 the absence of a geometrical vignetting. the effectivearea becomes Comparisor of Eqs. (17) ," In the absence of a geometrical vignetting, the effectivearea becomes Comparison of Eqs. \ref{eq:case1}) )"
and (18)) indicates that the effective area can be written as which represents the general expression for the effective area seen by à source on-axis., and \ref{eq:case2}) ) indicates that the effective area can be written as which represents the general expression for the effective area seen by a source on-axis.
" We can now compute the effective area for a source off-axis,", We can now compute the effective area for a source off-axis.
 Because of the axial symmetry of the mirror. we choose the X axis so that the source Hes in the vz plane. on the side of the positive x axis (refer to Fig. 5)).," Because of the axial symmetry of the mirror, we choose the $x$ axis so that the source lies in the $xz$ plane, on the side of the positive $x$ axis (refer to Fig. \ref{fig:mirror_scheme}) )."
 We define 8>0 to be the angle between z and the source direction., We define $\theta > 0$ to be the angle between $z$ and the source direction.
 With respect to the on-axis case. there are some additional difficulties.," With respect to the on-axis case, there are some additional difficulties."
" The ray no longer lies within a meridional plane of the mirror. so the polar angles of the impact positionson the primary and secondary mirror segment. q, and y>. differ n general. and the incidence angles ay and a> also vary with them."," The ray no longer lies within a meridional plane of the mirror, so the polar angles of the impact positionson the primary and secondary mirror segment, $\varphi_1$ and $\varphi_2$ , differ in general, and the incidence angles $\alpha_1$ and $\alpha_2$ also vary with them."
" Nevertheless. since the maximum distance of the two Npact points is ~Ly,+£2= 2L. the off-plane linear ""Sisplacement is 201, at most: therefore. στ-ϕι< 3901. which is in general negligible with respect to y,; and yo "," Nevertheless, since the maximum distance of the two impact points is $\sim L_1+L_2 \approx 2L$ , the off-plane linear displacement is $2\theta L$ at most: therefore, $|\varphi_2-\varphi_1|\lesssim 4\theta L'$ , which is in general negligible with respect to $\varphi_1$ and $\varphi_2$ "
bridge a gap of our understanding between (he electron injection and acceleration at the interplanetary/planetary bow shocks with relatively low Mach numbers ancl (he astrophysical shocks with very high Mach numbers.,bridge a gap of our understanding between the electron injection and acceleration at the interplanetary/planetary bow shocks with relatively low Mach numbers and the astrophysical shocks with very high Mach numbers.
 The authors are grateful to T. Terasaava for fruitful discussion., The authors are grateful to T. Terasawa for fruitful discussion.
 MII acknowledges support from the International Space Science Institute (ISSI) at Dern/Switzerlauxd for the collisionless shock working group., MH acknowledges support from the International Space Science Institute (ISSI) at Bern/Switzerland for the collisionless shock working group.
ou graphene. we propose that this polycrystalline model can solve the standing problem of the IS UW feature width by diut of its specific structure.,"on graphene, we propose that this polycrystalline model can solve the long-standing problem of the IS UV feature width by dint of its specific structure."
 The PC model was inspired by reflection nieasurements on industrial polycrystalline eraplite., The PG model was inspired by reflection measurements on industrial polycrystalline graphite.
 This is a hydrogeu-free asseiibly. of randomly oriented microscopic sp? carbon chips (or bricks). forming macroscopically homogeneous aud isotropic. solid. exaius usually delivered in the form of powders of different erauulomoetry. after mulling to different extents.," This is a hydrogen-free assembly of randomly oriented microscopic $^{2}$ carbon chips (or bricks), forming macroscopically homogeneous and isotropic, solid, grains usually delivered in the form of powders of different granulometry, after milling to different extents."
 In this form. the constituent eraphitic clips are usually io smaller than a few teus of nanometers in all three dimensions.," In this form, the constituent graphitic chips are usually no smaller than a few tens of nanometers in all three dimensions."
 They are pressed under more than 10 kbars iuto cisc-shaped pellets., They are pressed under more than 10 kbars into disc-shaped pellets.
 The spectral reflectance RCA) is usually measured at near normal incidence on the circular face of the disc. preferably after some simface smoothing.," The spectral reflectance $R(\lambda)$ is usually measured at near normal incidence on the circular face of the disc, preferably after some surface smoothing."
 The IaunuersdR&ronig (1) relations (see Bohlen aud ITuffinan (1983))) are then applied to R to deduce the bulk dielectric fuuctions of this material., The Kramers-Kronig (K-K) relations (see Bohren and Huffman \cite{boh}) ) are then applied to $R$ to deduce the bulk dielectric functions of this material.
 These are isotropic because. even under the applied pressure. the chips retain their random orientations.," These are isotropic because, even under the applied pressure, the chips retain their random orientations."
 These experimental fictions cau be compared with those obtained by applviug the Brugecinan mixing formula (see Bohlen aud Uniffinan (1983))) to a mixture. iu the ratio 1/3 to 2/3. of the dielectric functions of pure. bulk. graphite for E|c aud ELe orientations. respectively. as determined experimentally.," These experimental functions can be compared with those obtained by applying the Bruggeman mixing formula (see Bohren and Huffman \cite{boh}) ) to a mixture, in the ratio 1/3 to 2/3, of the dielectric functions of pure, bulk, graphite for $E\parallel c$ and $E\perp c$ orientations, respectively, as determined experimentally."
 The agreement is adequate provided the powder material coutaius no volatile atoms such as those that are carried bv the binding substances which are sometimes introduced for industrial purposes., The agreement is adequate provided the powder material contains no volatile atoms such as those that are carried by the binding substances which are sometimes introduced for industrial purposes.
" When the PC dielectric fictions are used to deduce the extinction eficiency of Ravleigli-sized erains. this is found to have a Fróllich (or surface. or pleuuon) resonance. nearly Loveutzian in profile. peaking ucar L6 yan ο. with a width FWHAM~L342n. ο,"," When the PG dielectric functions are used to deduce the extinction efficiency of Rayleigh-sized grains, this is found to have a Fröllich (or surface, or plasmon) resonance, nearly Lorentzian in profile, peaking near 4.6 $\mu$ $^{-1}$, with a width $\sim1.3 \mu$ $^{-1}$."
 Note that this width falls at the higher cud of the spectral interval within which most of the interstellar (IS) feature widths are observed (Fitzpatrick aud Massa (1986).. Fitzpatrick (2007)3).," Note that this width falls at the higher end of the spectral interval within which most of the interstellar (IS) feature widths are observed (Fitzpatrick and Massa \cite{fm}, Fitzpatrick \cite{fit}) )."
 There is no wav of reducing this width. let alone tailoring it. if the bulls graphite properties are retained.," There is no way of reducing this width, let alone tailoring it, if the bulk graphite properties are retained."
 Uvdrogenation. racdiation-induced defects. eran agelomerationao are all known to broaden aud weaken the plasimon feature.," Hydrogenation, radiation-induced defects, grain agglomeration are all known to broaden and weaken the plasmon feature."
 Ou the other haud. the receut fhary of research on the clectronic properties of eraplene provides a treasure trove of experimental facts related to our problem (see Cel aud Novoselov (20073)).," On the other hand, the recent flurry of research on the electronic properties of graphene provides a treasure trove of experimental facts related to our problem (see Geim and Novoselov \cite{gn}) )."
 Graphlieue is a sheet of compactly assembled hexagonal carbon riugs. which can uow be produced in the laboratory. or on a substrate. in uanometric sizes.," Graphene is a sheet of compactly assembled hexagonal carbon rings, which can now be produced in the laboratory, free-standing or on a substrate, in nanometric sizes."
 Multiple lavers can also be stacked upon one another in au orderly manner. as in natural eraphite (the so-called ADAD. or Bernal. stacking). or in so-called disorder. ic.," Multiple layers can also be stacked upon one another in an orderly manner, as in natural graphite (the so-called ABAB, or Bernal, stacking), or in so-called disorder, i.e."
We consider subsets of the full data set. to. construct our statistic comprising Nau data. points.,We consider subsets of the full data set to construct our statistic comprising $N_{\rm subset}$ data points.
 For analytic marginalisation we keep two anchor SNe aside., For analytic marginalisation we keep two anchor SNe aside.
 Since the ACDAL model fits the gold. data sets CD04 anc CGIDOT well. we first obtain the best. fit to the using the mareinalised likelihood function. given in Eq X5.. and then for cach SN we calculate the residualsO31)].. which is [ree of {ο by. virtue of Eq 4..," Since the $\Lambda$ CDM model fits the gold data sets GD04 and GD07 well, we first obtain the best fit to the using the marginalised likelihood function given in Eq \ref{eq:marglik}, and then for each SN we calculate the residuals, which is free of $H_0$ by virtue of Eq \ref{eq:diffmu}."
" Phe standard. error [or a SN at redshift z; is o,(2;). and we assume that the errors on SNe are statistically uncorrelated."," The standard error for a SN at redshift $z_i$ is $\sigma_{\mu}(z_i)$, and we assume that the errors on SNe are statistically uncorrelated."
 We define X5=\G)/ Nous where the marginalised Vay is defined in Eq. AT.," We define $\chi^2_R = \chi^2_M/N_{\rm
subset}$ , where the marginalised $\chi^2_M$ is defined in Eq. \ref{eq:mchisq}."
 v5 indicates the statistical scatter of the subset from the best fit I CDM mocel., $\chi^2_R$ indicates the statistical scatter of the subset from the best fit $\Lambda$ CDM model.
 Hs expectation value is unity (see Appendix for a proof). that is AA)= 1.," Its expectation value is unity (see Appendix for a proof), that is $\langle \chi^2_R
\rangle = 1$ ."
" We divide the data into two hemispheres labeled by the direction. vector P. ancl take the cillercnce of the X2, computed for the two hemispheres separately to obtain AZ=Vp,Ne where label 1 corresponds to that hemisphere towards which the direction vector P points and label 727 refers to the opposite hemisphere."," We divide the data into two hemispheres labeled by the direction vector $\hat{n}$, and take the difference of the $\chi_R^2$ computed for the two hemispheres separately to obtain $\Delta
\chi_{\hat{n}}^2 =\chi^2_{R1} - \chi^2_{R2} $, where label '1' corresponds to that hemisphere towards which the direction vector $\hat{n}$ points and label '2' refers to the opposite hemisphere."
 We take the absolute value of 42. since we are interested in the largest magnitude of this quantity.," We take the absolute value of $\Delta
\chi_{\hat{n_i}}^2$ since we are interested in the largest magnitude of this quantity."
 We then vary the direction n across the sky to obtain the maximum absolute difference As shown in GSLOS. the distribution of A\2 follows a simple. two parameter Gumbel distribution. characteristic of extreme value distribution tvpe LE (Ixendall&Stuart.1977:Cumbel1965 ).. where the position parameter m and the scale parameter s completely determine. the distribution.," We then vary the direction $\hat{n}$ across the sky to obtain the maximum absolute difference As shown in GSL08, the distribution of $\Delta_{\chi^2}$ follows a simple, two parameter Gumbel distribution, characteristic of extreme value distribution type I \citep{ken77, gum65}, where the position parameter $m$ and the scale parameter $s$ completely determine the distribution."
 Lo quantify departures from isotropy we need to know the theoretical clistribution. which is caleulated numerically by simulating several sets of Gaussian. distributed: y; on the gold. set SN positions and obtaining from cach realization.," To quantify departures from isotropy we need to know the theoretical distribution, which is calculated numerically by simulating several sets of Gaussian distributed $\chi_i$ on the gold set SN positions and obtaining $\Delta_{\chi^2}$ from each realization."
" For comparison with theory we 3,2follow GSLOS and. compute a bootstrap distribution by shullling the data values z;. σε) and eo,(2;) over the SNe positions (for further details. see 5105)."," For comparison with theory we follow GSL08 and compute a bootstrap distribution by shuffling the data values $z_i$, $\mu(z_i)$ and $\sigma_{\mu}(z_i)$ over the SNe positions (for further details see GSL08)."
 As mentioned in GSLOS. the above procedure separates the data sets into hot and cold SNe that have large and small dispersions with respect to the best fit model.," As mentioned in GSL08, the above procedure separates the data sets into hot and cold SNe that have large and small dispersions with respect to the best fit model."
 However. note that these two sets could still indicate the same cosmology. albeit with a different value of 47.," However, note that these two sets could still indicate the same cosmology, albeit with a different value of $\chi^2$."
 To. anicliorate this deficiency we now introduce a new statistic that does not suller from this artifact., To ameliorate this deficiency we now introduce a new statistic that does not suffer from this artifact.
 As mentioned above. x72 does not contain information about whether the SN is above or below the fit.," As mentioned above, $\chi_i^2$ does not contain information about whether the SN is above or below the fit."
 An obvious eencralization that does contain information regarding whether the SN at a redshift is closer or farther from us can be obtained by considering a statistic based on y;s. We consider two subsets of data defined. by. two hemispheres labeled by the direction. vector 5. containing μαι and Non λος where the total number of SNe. Noun and define the quantity Clearly (Ay)=0 and £GNY)521.," An obvious generalization that does contain information regarding whether the SN at a redshift is closer or farther from us can be obtained by considering a statistic based on $\chi_i$ s. We consider two subsets of data defined by two hemispheres labeled by the direction vector $\hat{n}$, containing $N_{\rm north}$ and $N_{\rm south}$ SNe, where the total number of SNe, $N = N_{\rm north} + N_{\rm south}$ , and define the quantity Clearly $\langle \Delta \chi_{\hat{n}} \rangle = 0$ and $\langle
(\Delta \chi_{\hat{n}})^2 \rangle = 1$."
 From the central limit theorem (Ixendall&Stuart1977). it follows that for ANOld. the quantity Ay follows a Caussian distribution with a zero mean and unit variance.," From the central limit theorem \citep{ken77} it follows that for $N \gg 1$, the quantity $\Delta\chi$ follows a Gaussian distribution with a zero mean and unit variance."
" As in the previous case we maximize this quantity by varving the direction P across the sky to obtain the maximum absolute dillerence Unlike the 4,2 statistic this statistic is not marginalisced over the Llubble constant since our results show that nmarginalising over it instead. of using its best fit value has only a marginal effect on 2X2.", As in the previous case we maximize this quantity by varying the direction $\hat{n}$ across the sky to obtain the maximum absolute difference Unlike the $\Delta_{\chi^2}$ statistic this statistic is not marginalised over the Hubble constant since our results show that marginalising over it instead of using its best fit value has only a marginal effect on $\Delta_{\chi^2}$.
 Moreover. in the limit Vv2»1. and assuming a uniform sky coverage. we expect the two hemispheres to contain roughly an equal number of SNe.," Moreover, in the limit $N \gg 1$, and assuming a uniform sky coverage, we expect the two hemispheres to contain roughly an equal number of SNe."
 In this case it is clear that AY would depend only weakly on Hua., In this case it is clear that $\Delta_{\chi}$ would depend only weakly on $H_0$.
 This statistic dillers from the previous one in that the Ay statistic has a theoretical limit where the position and the shape parameters can be determined analytically., This statistic differs from the previous one in that the $\Delta_{\chi}$ statistic has a theoretical limit where the position and the shape parameters can be determined analytically.
 Giver Ny independent directions on the sky we are essentially determining the maximum of a sample of size My where the incliviclual numbers are drawn from a Gaussian distribution with a zero mean and unit variance., Given $N_d$ independent directions on the sky we are essentially determining the maximum of a sample of size $N_d$ where the individual numbers are drawn from a Gaussian distribution with a zero mean and unit variance.
" In the limit NM,Z»1 the parameters are given by (Haan&Ferreira.2006) where we have to additionally assume that the number of SNe N° l.since the distribution for y becomes Gaussian only in this limit."," In the limit $N_d \gg 1$ the parameters are given by \citep{haan06} where we have to additionally assume that the number of SNe $N\gg 1$, since the distribution for $\chi$ becomes Gaussian only in this limit."
 This is convenient since at least [or large data sets. which will be available in the future. a comparison with theory becomes simpler.," This is convenient since at least for large data sets, which will be available in the future, a comparison with theory becomes simpler."
 However. for a smaller number of SNe there is a possibility that not. all clirections. are independent. in fact. it is quite possible that two directions contain exactly same subsets in the two hemisphere.," However, for a smaller number of SNe there is a possibility that not all directions are independent, in fact, it is quite possible that two directions contain exactly same subsets in the two hemisphere."
 In this situation is is clear that the total independent. directions 1s a smaller number than Αν and thus theoretical distribution would. be rightward shifted ancl also more sharply peaked., In this situation is is clear that the total independent directions is a smaller number than $N_d$ and thus theoretical distribution would be rightward shifted and also more sharply peaked.
 For this reason we also calculate the bootstrap cistribution and the theoretical distribution in the same manner as for the previous statistic., For this reason we also calculate the bootstrap distribution and the theoretical distribution in the same manner as for the previous statistic.
 In GSLOS we discussed. a specific bias in the bootstrap distribution. showing that it is shifted slightly to the left of the theoretical distribution due to the fact that theoretical distribution is obtained by assuming ys to be Caussian random variates with a zero meanand unit variance.," In GSL08 we discussed a specific bias in the bootstrap distribution, showing that it is shifted slightly to the left of the theoretical distribution due to the fact that theoretical distribution is obtained by assuming $\chi_i$ s to be Gaussian random variates with a zero meanand unit variance."
 Theoretical yes are unbounded. however the bootstrap distribution is obtained by shullling through a of v; where the x;s are obviously bounded.," Theoretical $\chi_i$ s are unbounded, however the bootstrap distribution is obtained by shuffling through a of $\chi_i$ where the $\chi_i$ s are obviously bounded."
 Lt is clear that on the average this should produce slightly smaller values of i2 in comparisonto what one expects fron a, It is clear that on the average this should produce slightly smaller values of $\Delta_{\chi^2}$ in comparisonto what one expects from a
we discuss our selection of the transit surveys inchided in our analysis aud the normalization of individual survey results.,we discuss our selection of the transit surveys included in our analysis and the normalization of individual survey results.
 In refsec:results we present our combined upper Bits., In \\ref{sec:results} we present our combined upper limits.
 refsec:discussion ancl §refsec:conclusion are devoted to discussion and ou conchisions. respectively.," \\ref{sec:discussion} and \\ref{sec:conclusion} are devoted to discussion and our conclusions, respectively."
— The techuiques used to fluc planets in auy pliotomietric survey are broadly similar., The techniques used to find planets in any photometric survey are broadly similar.
 All surveys attempt to achieve laugh photometric accuracy with maximal temporal coverage over the lougest possible period of time., All surveys attempt to achieve high photometric accuracy with maximal temporal coverage over the longest possible period of time.
 The better the temporal coverage. the more sensitive a survey is to all transits. and the louger the duration of the survey. the larger the range of planetary periods that are detectable.," The better the temporal coverage, the more sensitive a survey is to all transits, and the longer the duration of the survey, the larger the range of planetary periods that are detectable."
 We refer the reader to Pepper&Gaudi(2005) and vouBraunetal.(2005). for extensive discussions of the factors that determine a survey's seusitivitv to ransits and the design of successful surveys., We refer the reader to \citet{pepper2005} and \citet{vonbraun2005} for extensive discussions of the factors that determine a survey's sensitivity to transits and the design of successful surveys.
 Ounce the data have been collected. oue produces ight curves aud searches for periodic variabilitv.," Once the data have been collected, one produces light curves and searches for periodic variability."
 The x-fittiug. least-squares (BLS) method iutroduced. by Ílovácsetal(2002) is à popular means to do so within the transit commnity., The box-fitting least-squares (BLS) method introduced by \citet{kovacs2002} is a popular means to do so within the transit community.
 The aleoritlin searches for seriodic. rectangular deviations from a flat elt curve. and is an objective. repeatable meaus of ideutifvius ransit candidates.," The algorithm searches for periodic, rectangular deviations from a flat light curve, and is an objective, repeatable means of identifying transit candidates."
" Conversely, some authors may choose o identify transits by eve."," Conversely, some authors may choose to identify transits by eye."
 Reeardless of the method used to identify photometric eclipses. the plauctary lature of convincing transit candidates must then be confirmed with radial velocity follow-up observations. since astrophysical false positives cau iminide the transit of a Jupiter-sized object.," Regardless of the method used to identify photometric eclipses, the planetary nature of convincing transit candidates must then be confirmed with radial velocity follow-up observations, since astrophysical false positives can mimic the transit of a Jupiter-sized object."
 Iu the case of a null result. some authors choose to carefully quantity their detection cfiicicucies aud place upper lanits on the frequency of short-period plaucts.," In the case of a null result, some authors choose to carefully quantify their detection efficiencies and place upper limits on the frequency of short-period planets."
 One typically injects a Luge umber of sinmlated. inb-darkened transits iuto coustant sinmlated stars or ight curves from the survey itself. aud then subjects τοσο transits to the same detection algorithms with re same selection criteria as the original data.," One typically injects a large number of simulated, limb-darkened transits into constant simulated stars or light curves from the survey itself, and then subjects these transits to the same detection algorithms with the same selection criteria as the original data."
 The raction of iujected transits that are recovered quautifies ιο detection efficiency., The fraction of injected transits that are recovered quantifies the detection efficiency.
 This efficiency. is a complicated conibination of effects that is discussed iun more detail iu 1e next section., This efficiency is a complicated combination of effects that is discussed in more detail in the next section.
 While the general process is simular for all survevs. ιο details of how cach author chooses to perform the steps varies slightly. aud in practice it is not. necessarily correct to directly compare or combine the results of wo surveys.," While the general process is similar for all surveys, the details of how each author chooses to perform the steps varies slightly, and in practice it is not necessarily correct to directly compare or combine the results of two surveys."
 Authors may make different asstmuptious about the period distribution of planets. or quote their results for cifferiue planetary radii etc.," Authors may make different assumptions about the period distribution of planets, or quote their results for differing planetary radii, etc.,"
 all of which are nuportant when comparing the results of multiple different surveys., all of which are important when comparing the results of multiple different surveys.
 In order to combine the results of several surveys. we lust re-nornalize. to the exteut that we can. fo a colon set of assuuptious.," In order to combine the results of several surveys, we must re-normalize, to the extent that we can, to a common set of assumptions."
 Iu this section we discuss the mathematical description of the detection efficicucy of a transit survey. the calculation of upper limits ou the frequency. of plaucts using a uull result. aud the method for combining several normalized surveys iuto a single upper limit.," In this section we discuss the mathematical description of the detection efficiency of a transit survey, the calculation of upper limits on the frequency of planets using a null result, and the method for combining several normalized surveys into a single upper limit."
" Tn general. the expected nuuber of planets detected iu the radius rauge Z2, to R,|dR, aud period range P to P|dP can bo written (Burkectal.2006) where & is the star. the sun is over all NY stars. and Prem. 1s the probability that the star is a cluster member."," In general, the expected number of planets detected in the radius range $R_p$ to $R_p +dR_p$ and period range $P$ to $P+dP$ can be written \citep{burke2006}
 where $k$ is the star, the sum is over all $N_{\star}$ stars, and $\mathcal{P}_{mem,k}$ is the probability that the star is a cluster member."
" f, is the fraction of stars that host short-period planets distributed as dPΠΙΟ.", $f_p$ is the fraction of stars that host short-period planets distributed as $d^2\mathcal{P}/dR_pdP$.
 We assmnue PPΠ is independent of the stellar ass., We assume $d^2\mathcal{P}/dR_pdP$ is independent of the stellar mass.
 This distribution is poorly constrained by exoplanuct statistics. and is usually assuned as a prior.," This distribution is poorly constrained by exoplanet statistics, and is usually assumed as a prior."
" We will assume that this distribution is a delta functiou in the radius aud nuiform in the οσαται of the period: where ""» is the planetary radius of interest.", We will assume that this distribution is a delta function in the radius and uniform in the logarithm of the period: where $R_p^{'}$ is the planetary radius of interest.
" P; is the geometric probability that a planet trausits its Lost star. such that where δὲ, is the radius of the star and « is the seniuajor axis of the orbit. where we have assumed the orbit is Pee Is the probability that a planet of radius δὲ, aud period P will be detected around star & if it transits. even the precise temporal coverage. precision. and signal-to-noise ratio of the observations."," $\mathcal{P}_{tr,k}$ is the geometric probability that a planet transits its host star, such that where $R_{\star}$ is the radius of the star and $a$ is the semimajor axis of the orbit, where we have assumed the orbit is $\mathcal{P}_{det,k}$ is the probability that a planet of radius $R_p$ and period $P$ will be detected around star $k$ if it transits, given the precise temporal coverage, precision, and signal-to-noise ratio of the observations."
 Iu practice Py. is trivial to calculate. while ως is cuphatically uot so.," In practice $\mathcal{P}_{tr}$ is trivial to calculate, while $\mathcal{P}_{det}$ is emphatically not so."
 It is iu general. the sinele most dificult factor to characterize in the cutive foriualisuu because if represents a couples interplay between nunerous different observational effects;," It is, in general, the single most difficult factor to characterize in the entire formalism because it represents a complex interplay between numerous different observational effects."
 Simple analytic estimates of P4 teud to overestimate the probability that the survey will be capable of detecting transits (Beatty&Candi 2008)., Simple analytic estimates of $\mathcal{P}_{det}$ tend to overestimate the probability that the survey will be capable of detecting transits \citep{beatty2008}.
. The best way to characterize Pg Is to performs Monte Carlo simulations with transits iujected iuto the elt curves of constant stars observed in the survey or simulated coustaut stars., The best way to characterize $\mathcal{P}_{det}$ is to perform Monte Carlo simulations with transits injected into the light curves of constant stars observed in the survey or simulated constant stars.
 When authors choose to perform an analysis of the detection efficiencies with Monte Carlo simulations. the resulting efüciencies generally account for P4. P. and dPdidP.," When authors choose to perform an analysis of the detection efficiencies with Monte Carlo simulations, the resulting efficiencies generally account for $\mathcal{P}_{det}$, $\mathcal{P}_{tr}$, and $d^2\mathcal{P}/dR_pdP$."
 Dy injecting trausits with raudouily siuupled paraiaueters from au assumed period distribution iuto stars in the survey that were searched for trausits. one automatically accounts for the iufrinsic ασ function of stars and planet period distributions.," By injecting transits with randomly sampled parameters from an assumed period distribution into stars in the survey that were searched for transits, one automatically accounts for the intrinsic mass function of stars and planet period distributions."
/ Factors such as noise present in the helt curves aud the observing window. which are crucial to determining Έρως. are also automatically included im transit jection and recovery schemes.," Factors such as noise present in the light curves and the observing window, which are crucial to determining $\mathcal{P}_{det}$, are also automatically included in transit injection and recovery schemes."
 ο Is generally determined separately through proper motion mecasurcuicuts. proximity of a star to the cluster main sequence. off-cluster comparison fields. or amodeling of the feld star population to determine contamination levels.," $\mathcal{P}_{mem}$ is generally determined separately through proper motion measurements, proximity of a star to the cluster main sequence, off-cluster comparison fields, or modeling of the field star population to determine contamination levels."
 After the detection cficicucies are quantified. one has," After the detection efficiencies are quantified, one has"
stavs. ducluding orbital circularization. sviTronizatiou between the orbital period and t1ο stars rotation. alc spin-orbit aliguinent (AMazeh2008).. have a time scidc5 that is strongly depeudeut on he scaled senijor axisM. (01Ry (e.g.Zahn1989).,"stars, including orbital circularization, synchronization between the orbital period and the stars' rotation, and spin-orbit alignment \citep{mazeh08}, have a time scale that is strongly dependent on the scaled semi-major axis, $a/R_p$ \cite[e.g.,][]{zahn89}."
. The spiu-orbit aliene spccifically has a relatively shor ine scale. two or thr orders of magnitude shorter thaw he circularization ti5 scale since it involves a transter of a smaller anit of angular momentum (Mazeh2008).," The spin-orbit alignment specifically has a relatively short time scale, two or three orders of magnitude shorter than the circularization time scale since it involves a transfer of a smaller amount of angular momentum \citep{mazeh08}."
. Therefore. is expected that for short perio nares the spi1 ANCS of the two stars had already aligned with the «ybial 1ionientuna axis. and anv mitial missaliguiucut are lic from their formation historv hacl already faclec aweuU.," Therefore, it is expected that for short period binaries the spin axes of the two stars had already aligned with the orbital momentum axis, and any initial missalignment — a relic from their formation history — had already faded away."
 Since for longer period svstcis. with a wider orbit. this process is slower. they are better targets for studyue the initia aligninent aud the svstenis: formatio1 process.," Since for longer period systems, with a wider orbit, this process is slower, they are better targets for studying the initial alignment and the systems' formation process."
" For exije. the missaligued svsteii DI Wer (Albrechtet 2009.. see also Mazel 2008)) has au orbital iod of 10.6 clay,"," For example, the missaligned system DI Her \citealt{albrecht09}, see also \citealt{mazeh08}) ) has an orbital period of 10.6 day."
 We used the binary caalog of Prsaetal. (2011).. including|1879 eclipsingsvstenis. to estimate the uunuber of eclipsing binary svstenis whose photometry will be accurate enough to allow looking or the PRM signal.," We used the eclipsing binary catalog of \cite{prsa11}, including 1879 systems, to estimate the number of eclipsing binary systems whose photometry will be accurate enough to allow looking for the PRM signal."
 We looked for svstenis whose eclipse duration 1s longero than 5 hours. and that eiveu their orvital period and brightness magnitude) the phoomoetrie S/N will be no less than half that of the 10 davs aud 12th mae svsteni examined here.," We looked for systems whose eclipse duration is longer than 5 hours, and that given their orbital period and brightness magnitude) the photometric S/N will be no less than half that of the 10 days and 12th mag system examined here."
 This includes 10 day systenis whose brightuess is down to 13.5 Thismae. or 12 mae systems with period up to 10 davs.," This includes 10 day systems whose brightness is down to 13.5 mag, or 12 mag systems with period up to 40 days."
 cut left about 320 systems., This cut left about 320 systems.
 When we restrict ourselves to low lass ratio svstenis. simular to the one we exanuned here. this uunber is about 50.," When we restrict ourselves to low mass ratio systems, similar to the one we examined here, this number is about 50."
 A detection of the PRM effect aud a measurement of the spin-orbit anele for the primary iu evel soue of those svstenis will ercatly extend the sample of eclipsing binaries for which tjs angle is kuown., A detection of the PRM effect and a measurement of the spin-orbit angle for the primary in even some of those systems will greatly extend the sample of eclipsing binaries for which this angle is known.
 Ta our nuunerical analysis we ignored several low-auplitude effects. including elipsoidal distortion due tidal forces. betweei the two stars and reflection of ight. or heating of one star by the other.," In our numerical analysis we ignored several low-amplitude effects, including ellipsoidal distortion due to tidal forces between the two stars and reflection of light, or heating of one star by the other."
 For a 10 day period binary both these effects are expected to lave an amplitude along the orbit which is smaller han that of the orbital |)'anunue effect (Zuckerctal. 2007)., For a 10 day period binary both these effects are expected to have an amplitude along the orbit which is smaller than that of the orbital beaming effect \citep{zucker07}.
. Iun addition. during eclipses the ellipsoidal aud reflection effects are at extrema. ucaning their iuduced Hux variation is mininual. while the orbital beanuue fiux is af its phase of fastcst varlation," In addition, during eclipses the ellipsoidal and reflection effects are at extrema, meaning their induced flux variation is minimal, while the orbital beaming flux is at its phase of fastest variation."
 The light curves of he ellipsoidal aud. reflection effects divine eclipse are expected to be sviuuetric about the aid eclipse. time. and. as for the orbital beamine. hey can be predicted using modoeliug of the out of eclipse light curve (Faigler&Mazel 2011).," The light curves of the ellipsoidal and reflection effects during eclipse are expected to be symmetric about the mid eclipse time, and, as for the orbital beaming, they can be predicted using modeling of the out of eclipse light curve \citep{faigler11}."
. We note that we igjore here t1e possibility of othe sources of low level noisc. such as stellar noise. mea hat this measuremen requires quiet stars.," We note that we ignore here the possibility of other sources of low level noise, such as stellar noise, meaning that this measurement requires quiet stars."
" Iu case f eclipsed. star shows slight activity in the form of Sl spots on its stwhace. this can lead to sMusoklal-li""m Hux variability at the rotational period."," In case the eclipsed star shows slight activity in the form of small spots on its surface, this can lead to sinusoidal-like flux variability at the rotational period."
 However. fo he type of svseni simulated here the orbital aud x rotation periods are not necessarily svuchronized. so th spots-iucdiced fux variability will be averaged ott oc olded on he orvital period. which is determined o hig accuracy |w the measurement of eclipses timine," However, for the type of system simulated here the orbital and self rotation periods are not necessarily synchronized, so the spots-induced flux variability will be averaged out once folded on the orbital period, which is determined to high accuracy by the measurement of eclipses timing."
 Qur methacl jas several possible degeneracies. whic[um if not carefully considered can lead to incorrect interpretation of the light curve. aud a biased value for A. Significant eccentricity can produc| asviunietrie eclipse helt curves. dependingC» ou the argumentOo of periasTOL. w (see Lapping2XS. for a description of the impact of eccentricity on pl:uietary transit lielit curves).," Our method has several possible degeneracies, which if not carefully considered can lead to incorrect interpretation of the light curve, and a biased value for $\lambda$ Significant eccentricity can produce asymmetric eclipse light curves, depending on the argument of periastron, $\omega$ (see \citealt{kipping08} for a description of the impact of eccentricity on planetary transit light curves)."
 Although. eccentricity can |)o estimated from the orbital beaming helt curve. aud. the value of ecosw Is accurately determined from he time difference between the primary and secondary οςIpses.," Although, eccentricity can be estimated from the orbital beaming light curve, and, the value of $e\cos \omega$ is accurately determined from the time difference between the primary and secondary eclipses."
 A second deegcucracy results from the PRAL signal beige nünücked by a sinall shift in iid eclipse time. which may eive the false impression that the PRM signal does not exist in the Πο curve. sugeesting a scenario where ;/ = 90 deg and A = 90 dee. shown in left panel.," A second degeneracy results from the PRM signal being mimicked by a small shift in mid eclipse time, which may give the false impression that the PRM signal does not exist in the light curve, suggesting a scenario where $i$ = 90 deg and $\lambda$ = 90 deg, shown in left panel."
 We looked into this possibility more closely ly Sbtracting the eclipse ligbit curve inchiding a simall shift iu mud eclipse time. aud without the PRM effect. from the lie;ht curve inchiding the PRM signal aud no time shift.," We looked into this possibility more closely by subtracting the eclipse light curve including a small shift in mid eclipse time, and without the PRM effect, from the light curve including the PRM signal and no time shift."
 The aid eclipse time shifts we examined were up to a few seconds. im sub second steps.," The mid eclipse time shifts we examined were up to a few seconds, in sub second steps."
 We found that the peals to peak amplitudes of the difference Πο curves were at least GO ppim. and the smallest difference is for time shifts close to Ls. We expect photometry obtained during the time span discussed here to allow aid eclipse measurements of better than 1 s precision.," We found that the peak to peak amplitudes of the difference light curves were at least 60 ppm, and the smallest difference is for time shifts close to 1 s. We expect photometry obtained during the time span discussed here to allow mid eclipse measurements of better than 1 s precision."
 Therefore. a solution which docs not include the PRAL effect (or assumes / = 90 deg aud A = 90 deg) i e a result eives a wrong mid eclipse time due to the sabove mentioned degeneracy will have a sguificautly lower likelihood than a solution that does include the PRM cetfect.," Therefore, a solution which does not include the PRM effect (or assumes $i$ = 90 deg and $\lambda$ = 90 deg) and as a result gives a wrong mid eclipse time due to the above mentioned degeneracy will have a significantly lower likelihood than a solution that does include the PRM effect."
 Modelue the PRM effect adds two parameters to the model. Vsin/ aud A of the eclipsed star.," Modeling the PRM effect adds two parameters to the model, $V_{rot}\sin I$ and $\lambda$ of the eclipsed star."
 Aw independent estinate of V;inZ can be achieved from a spectra of the star. thereby coustrainiug this parameter so the information in the PRM signal is used for fitting oulv a single parameter (A).," An independent estimate of $V_{rot}\sin I$ can be achieved from a spectrum of the star, thereby constraining this parameter so the information in the PRM signal is used for fitting only a single parameter $\lambda$ )."
" Moreover. au incepenudeut estimate of Vi,,s1nJ is müportaut for lifting a third degeneracy m our model. between A and V,,;*n. for edee-on orbits (e...Gaudi&Winn2007)."," Moreover, an independent estimate of $V_{rot}\sin I$ is important for lifting a third degeneracy in our model, between $\lambda$ and $V_{rot}\sin I$, for edge-on orbits \citep[e.g.,][]{gaudi07}."
. Another possible degencracy may result from the IuOrD signal being opposite iu sien to the PRM sigual (see Figures 1. aud 2)). eiviug a false impression that the PRM signal is samaller in amplitude.," Another possible degeneracy may result from the InOrB signal being opposite in sign to the PRM signal (see Figures \ref{fig:F8K0} and \ref{fig:lambda}) ), giving a false impression that the PRM signal is smaller in amplitude."
 In ecueral. for nou- aud non edge-on oricutations. the combined leh curve shape could be differeut than the PRM liebt curve. possibly leading to a biased estimate of A. even though for the primary eclipse shown iu Figures 1 and 2 the PRA sienal is clearly dominant.," In general, for non-aligned and non edge-on orientations, the combined light curve shape could be different than the PRM light curve, possibly leading to a biased estimate of $\lambda$, even though for the primary eclipse shown in Figures \ref{fig:F8K0} and \ref{fig:lambda} the PRM signal is clearly dominant."
 However. as already noted. the IuO:D lieht curve shape is insensitive to the spiu-orbi anele or the stars rotation velocity.," However, as already noted, the InOrB light curve shape is insensitive to the spin-orbit angle or the star's rotation velocity."
 The only parameter to which both signals are scusitive to is the radii ratk) re which is accurately determined from the eclipse lel curve even without accounting for the beanine-induced effects.," The only parameter to which both signals are sensitive to is the radii ratio $r$, which is accurately determined from the eclipse light curve even without accounting for the beaming-induced effects."
 Therefore. the InOrD light curve is determined independently of the PRM light curve. eliniuatiug this possible degeneracy.," Therefore, the InOrB light curve is determined independently of the PRM light curve, eliminating this possible degeneracy."
 Lastly. we note that the oblateness of stars. due to self rotation. nav also affect eclipse light curves. especially when oblateness leads to significant eravity darkening.," Lastly, we note that the oblateness of stars, due to self rotation, may also affect eclipse light curves, especially when oblateness leads to significant gravity darkening."
 Barues(2009) showed how this wil distort the shape of planetary transit leh curves when the host star is a rapidly rotating carly type star., \cite{barnes09} showed how this will distort the shape of planetary transit light curves when the host star is a rapidly rotating early type star.
However. since Sun-like stars are slow rotators this is expec‘ted to lave a miuor,"However, since Sun-like stars are slow rotators this is expected to have a minor"
ccombined data sets.,combined data sets.
 These profiles are displaved in Figure 3., These profiles are displayed in Figure 3.
 They are well characterized by a single. broad peaked pulse witha ~50% duty evele.," They are well characterized by a single, broad peaked pulse with a $\sim 50\%$ duty cycle."
 The profiles appear similar., The profiles appear similar.
 The pulsed cinission comprises 5.5% of the total ccouuts in the backerounc-subtracted folded light curves. sugecsting that steady emission from a svuchrotron nebula nuelt be dominating the fiux.," The pulsed emission comprises $5.5\%$ of the total counts in the background-subtracted folded light curves, suggesting that steady emission from a synchrotron nebula might be dominating the flux."
 This is the same fraction as measured for the Crab pulsar (Seward 1951)., This is the same fraction as measured for the Crab pulsar (Seward 1984).
 The sspectit of Ies 75 with its hard wou-thermal componcut is presented in Blanton elfaud (2000)., The spectrum of Kes 75 with its hard non-thermal component is presented in Blanton Helfand (2000).
 We have analyzed the sspectium of the system above 3 keV so as to exclude remmaut’s thermal emüssiou., We have analyzed the spectrum of the system above $3$ keV so as to exclude remnant's thermal emission.
 The 20 keV sspectrun was fitted to an absorbed power Luv model with an absorbing column of Njj=3.1«1077 7. the vvalue. aud used the standard background model along with a Gaussian dune at 6.5 keV corresponding to the Galactic Ridge diffuse Fe cinission.," The $3 - 20$ keV spectrum was fitted to an absorbed power law model with an absorbing column of $_H = 3.1 \times
10^{22}$ $^{-2}$, the value, and used the standard background model along with a Gaussian line at 6.5 keV corresponding to the Galactic Ridge diffuse Fe emission."
 We obtained a sood fit for a photon iudex P=2.1840.01 (\2= 1.1). typical of a voung. Crab-like pulsar and consistent with the vvalue.," We obtained a good fit for a photon index $\Gamma = 2.18 \pm\ 0.04$ $\chi_{\nu}^2 = 1.4$ ), typical of a young, Crab-like pulsar and consistent with the value."
" The unabsorbed flux from the pulsar plus svuchrotron nebula cuuission is L8+0.1&10 ere cin2s 1,", The unabsorbed flux from the pulsar plus synchrotron nebula emission is $1.8 \pm\ 0.4 \times 10^{-11}$ erg $^{-2}$ $^{-1}$.
 Next. we considered the spectra of the pulsed cluission alone by analyzing phase cepeudeut spectra.," Next, we considered the spectrum of the pulsed emission alone by analyzing phase dependent spectra."
 Using the off-pulse spectrum as background. we measure a hard photon iudex D—Llc0.3 (AZ= 1.5) with au uuabsorbed flux of the pulsed component of 0.96—.10P7» ewe 57s +t in: the 3.0 keV.| enerev baud.," Using the off-pulse spectrum as background, we measure a hard photon index $\Gamma = 1.1 \pm\ 0.3$ $\chi_{\nu}^2
= 1.5$ ) with an unabsorbed flux of the pulsed component of $0.96 \pm\
0.2 \times 10^{-12}$ erg $^{-2}$ $^{-1}$ in the $3-10$ keV energy band."
" The steady increase in the period suggests a coustaut enerev loss at a rateofE=(23?IP/D?~10""Teressd. where Z ds the NS moment of inertia in units of ο en."," The steady increase in the period suggests a constant energy loss at a rate of ${\dot E} = (2\pi)^2 I{\dot P}/P^{3} \simeq 10^{37} I {\rm \
ergs \ s^{-1}}$, where $I$ is the NS moment of inertia in units of $1.4 \times 10^{45}$ g $^{2}$."
" In the standard pulsar model of a rotating magnetized dipole. spin-down οσον is lost via magnetic dipole vacation. E«Pτος(68),"," In the standard pulsar model of a rotating magnetized dipole, spin-down energy is lost via magnetic dipole radiation, ${\dot E} \sim (B_p^2 R^6 \Omega^4)/(6c^3)$."
 This model assunies a braking index of 3 and a uniformly maenuetized stellar iuterior., This model assumes a braking index of 3 and a uniformly magnetized stellar interior.
" From the above relatiouships we can inter the surface magnetic Ποια at the pole. which. for0258. gives D,cL8«1075 C. Tere. R~ is the neutron star radius. O=2z/P is the augular velocity of the rotation. and ο is the speed of lightvacuo."," From the above relationships we can infer the surface magnetic field at the pole, which, for, gives $B_p \simeq
4.8\times 10^{13}$ G. Here, $R\sim 10{\rm\ km}$ is the neutron star radius, $\Omega=2\pi/P$ is the angular velocity of the rotation, and $c$ is the speed of light."
 The characteristic spindown age (7=P2P) of lis 2723 vr. the vouugest characteristic age of any known pulsar.," The characteristic spindown age $\tau=P/{2\dot{P}}$ ) of is $\simeq 723$ yr, the youngest characteristic age of any known pulsar."
 For a well-behaved Crab-like pulsar with a braking iudex siuiar to 3. the spincdown age is reasonably consisteut with that reported for Nes 75 itself.," For a well-behaved Crab-like pulsar with a braking index similar to 3, the spindown age is reasonably consistent with that reported for Kes 75 itself."
 Blanton Ilelfaud (1996) estimated the age of the remuaut to lie between 900 aud 1.300 vears. based on free expansion aud Sedov phase estimates. respectively.," Blanton Helfand (1996) estimated the age of the remnant to lie between 900 and $4,300$ years, based on free expansion and Sedov phase estimates, respectively."
 They preferred a value closer o 10 vears. while aveuing that the remmant was still in Yoo expansion phase. giving an age r9004190-lovears.," They preferred a value closer to $10^3$ years, while arguing that the remnant was still in free expansion phase, giving an age $\tau \sim 900 d_{19}v_5^{-1}$ years."
 Tere. the distance to the SNR is 194454 kpc (Milue 1979) and the free expansion velocity is 500005 lau +.," Here, the distance to the SNR is $19d_{19}$ kpc (Milne 1979) and the free expansion velocity is $5000v_5$ km $^{-1}$."
 Cüveu he quality of age estimators. the two nuits are cousisteut.," Given the quality of age estimators, the two limits are consistent."
 Finally. we note that the cumulative enerev in particles and fields within the plerion a calorimetric measure of he pulsars activity is estimated frou radio observations o he ]DM“ore (Blauton Ποια 1996).," Finally, we note that the cumulative energy in particles and fields within the plerion – a calorimetric measure of the pulsar's activity – is estimated from radio observations to be $10^{48}d_{19}^{17/7}$ erg (Blanton Helfand 1996)."
 This is also consistent with the lower limit derived from the age of the oulsar and its current spin-cown loss. rE3.1057 cre.," This is also consistent with the lower limit derived from the age of the pulsar and its current spin-down loss, $\tau \dot E \sim 3\times 10^{47}$ erg."
 From the spin properties aud location of wwe infer an extremely voune and hishlv magnetized oulsar associated with the SNR Ίος 75., From the spin properties and location of we infer an extremely young and highly magnetized pulsar associated with the SNR Kes 75.
" At the assmmed distance to the SNR. the total N-rayv tuinosity from the pulsar plus svuchrotron nebula is TASLOMdT, eves linthe3/ 10 keV band: for comparison. his translates to 2.1«1076 eres tin the ROSATO.1 2.1 τον. baud for the spectral parameters eiven above."," At the assumed distance to the SNR, the total X-ray luminosity from the pulsar plus synchrotron nebula is $7.8 \times 10^{35}d_{19}^2$ erg $^{-1}$ in the $3-10$ keV band; for comparison, this translates to $2.1 \times 10^{36}$ erg $^{-1}$ in the ROSAT $0.1-2.4$ keV band for the spectral parameters given above."
 The N-rav luminosity of this region was already found to be he amoueg the highest of auv of the Crab-like SNRs. and second only to the Crab itself. depending on the true distance.," The X-ray luminosity of this region was already found to be the among the highest of any of the Crab-like SNRs, and second only to the Crab itself, depending on the true distance."
" The derived pulsed luminosity of L1«Loe, ere + (3.10 keV) suggests that ~1«10? of the oulsaurs spin-down energy is enadtted as pulsed N-ravs iu he 3.10 keV band or —1.10° in the Q124 τον. This value is simular to those observed from other rotation powered pulsars (Becker Titimuuper 1999) aud suggests that particle acceleration is occurming iu the NS uaenetosphere."," The derived pulsed luminosity of $4.1 \times 10^{34}
d_{19}^2$ erg $^{-1}$ $3-10$ keV) suggests that $\sim 4 \times
10^{-3}$ of the pulsars spin-down energy is emitted as pulsed X-rays in the $3-10$ keV band or $\sim 1 \times 10^{-3}$ in the $0.1-2.4$ keV. This value is similar to those observed from other rotation powered pulsars (Becker Trümmper 1999) and suggests that particle acceleration is occurring in the NS magnetosphere."
" However. the ratio of the total pulsar plus icbula luminosity to the spin-down energy loss is GZ, ines ercater than that for the Crab pulsar."," However, the ratio of the total pulsar plus nebula luminosity to the spin-down energy loss is $d_{19}^2$ times greater than that for the Crab pulsar."
 Tn many wavs. ireselnbles auv other rotation-powered Crab-like pulsar: jowever. its period. spiu-down rate. spiu-down conversion efficiency. and inferred maguetie field. are cach au orderofiuaguitude ereater.," In many ways, resembles any other rotation-powered Crab-like pulsar; however, its period, spin-down rate, spin-down conversion efficiency, and inferred magnetic field, are each an order-of-magnitude greater."
" The timing paramcters of the lew pulsar are imnost similar to the voune 0. bs pulsar J1119 G127 with its period derivative of LO.1012 s/s nupbiug B,=LAs1015 C. This pulsu. though. { times closer. does not coutain a simular bright radio or XN-rav pleriou (Camilo ot al."," The timing parameters of the new pulsar are most similar to the young $0.4$ -s pulsar $-$ 6127 with its period derivative of $4.0 \times 10^{-12}$ s/s implying $B_p = 4.1
\times 10^{13}$ G. This pulsar, though 4 times closer, does not contain a similar bright radio or X-ray plerion (Camilo et al."
 2000)., 2000).
 Recent research suggests that vouug ueutron stars niv have at least two distinct evolutionary branches (see CGottholf Vasisht 2000 aud refs., Recent research suggests that young neutron stars may have at least two distinct evolutionary branches (see Gotthelf Vasisht 2000 and refs.
 therein)., therein).
 Besides the Crab-like pulsus which evolve through magnetic braking. with fields in the range 1012.—1079 Cb and spin periods," Besides the Crab-like pulsars which evolve through magnetic braking, with fields in the range $10^{12} - 10^{13}$ G and spin periods"
"As shown in the previous sections, the photospheric velocities of four SNe in our sample evolved similarly.","As shown in the previous sections, the photospheric velocities of four SNe in our sample evolved similarly."
" SNe 1999em, 2004et and 2006bp had high velocities at early phases and they decreased quickly, although their decline slopes were different."," SNe 1999em, 2004et and 2006bp had high velocities at early phases and they decreased quickly, although their decline slopes were different."
" SN 2004dj probably showed similar evolution, but the lack of the early-phase data prevents a more detailed comparison."," SN 2004dj probably showed similar evolution, but the lack of the early-phase data prevents a more detailed comparison."
" On the contrary, SN 2005cs was a very different, low-energy SN II-P as discussed in detail in previous studies."," On the contrary, SN 2005cs was a very different, low-energy SN II-P as discussed in detail in previous studies."
 It had lower early velocities and the velocity curve decreased much faster than for all the other SNe., It had lower early velocities and the velocity curve decreased much faster than for all the other SNe.
Fiewre 10. displays profiles for the axial field case with Ey=0.5 at a fixed value of sin?;=0.1 but with differcut field values at the iucr disk radius of w=1. of loghy=2.1.0.12.|E to achieve a large dynamic ranec iu Tanle ratios throughout the disk.,"Figure \ref{figApp2} displays profiles for the axial field case with $E_1=0.5$ at a fixed value of $\sin^2 i = 0.4$ but with different field values at the inner disk radius of $\varpi=1$, of $\log b_0
= -2, -1, 0, +2, +4$ to achieve a large dynamic range in Hanle ratios throughout the disk."
" As by increases. the location where B—Dy, moves outward to ze=by."," As $b_0$ increases, the location where $B=\BHan$ moves outward to $\varpi = b_0$."
 The upper left paucl in Figure δ shows normalized profiles of FP: lower paucls slow the fractional polarizations aand45:5 auc upper right shows the sshapes.," The upper left panel in Figure \ref{fig8} shows normalized profiles of ${\cal F}_I^{\rm
sc}$; lower panels show the fractional polarizations and; and upper right shows the shapes."
 The color sequencing is the same as in Figure 5.., The color sequencing is the same as in Figure \ref{fig5}.
 In laree part the effect of an axial field is to rotate and foreshorten the Q-U segments relative to the zero field case., In large part the effect of an axial field is to rotate and foreshorten the Q-U segments relative to the zero field case.
 For a toroidal maeuetic field with B=£z. the spherical geometry for the scattering problems. is moderately complex.," For a toroidal magnetic field with $\vec{B}=B_\varphi\,\hat{\varphi}$, the spherical geometry for the scattering problem is moderately complex."
 It is not possible to write down simple complete expressions for the phase scattering functions., It is not possible to write down simple complete expressions for the phase scattering functions.
 But as demonstrated in Paper IL. there are still some special cases that are analytic.," But as demonstrated in Paper II, there are still some special cases that are analytic."
 For example in the saturated limit. FPS=0. aud the Γ and Q fluxes become For a disk viewed edee-on. solutions for the Stokes fluxes cannot be derived analytically: however. the scattering functions simply to Examples of scattering aud polarized profiles are shown in Figure[m] LL as well as Q-U plots across the polarized profiles.," For example in the saturated limit, ${\cal F}_U^{\rm sc} = 0$, and the $I$ and $Q$ fluxes become For a disk viewed edge-on, solutions for the Stokes fluxes cannot be derived analytically; however, the scattering functions simplify to Examples of scattering and polarized profiles are shown in Figure \ref{figApp3} as well as Q-U plots across the polarized profiles."
 Iu the previous[m] section. sine£=0.1 was used for au axial field to give a siguificaut signal in67.," In the previous section, $\sin^2 i=0.4$ was used for an axial field to give a significant signal in."
. For this toroidal field case. an edee-onc» disk with sint;=1.0 was used.," For this toroidal field case, an edge-on disk with $\sin^2 i=1.0$ was used."
 The style of this figurec» is the same as Figure[m] 10..., The style of this figure is the same as Figure \ref{figApp2}.
 The a iiajor distinctive in relatiou to a disk with au axial field is that a toroidal field leads to pprofiles that are autisvunmetric instead of svuuuctric., The a major distinctive in relation to a disk with an axial field is that a toroidal field leads to profiles that are antisymmetric instead of symmetric.
phasing up multiple ALMA antennas. the resulting constraint is strengthened dramatically. in. principle allowing à spin magnitude estimate with precision +0.1.,"phasing up multiple ALMA antennas, the resulting constraint is strengthened dramatically, in principle allowing a spin magnitude estimate with precision $\pm0.1$."
 The spin constraints. associated with observations with a phased ALMA are shown in Figure 4ee. at which point the allowed model parameter space becomes too small to resolve with the parameter sampling we employed.," The spin constraints associated with observations with a phased ALMA are shown in Figure \ref{fig:CPspinlimits}e e, at which point the allowed model parameter space becomes too small to resolve with the parameter sampling we employed."
 When the LMT is included (Figure 4ff). the allowed region is reduced to a single grid point in our model space (Ad=0.01. A0=&= 1°).," When the LMT is included (Figure \ref{fig:CPspinlimits}f f), the allowed region is reduced to a single grid point in our model space $\Delta a=0.01$, $\Delta\theta=\Delta\xi=1^\circ$ )."
 Whether such high precision spin. estimation is possible in practice is almost certainly going to be limited by the uncertainties in the modeling of the emission region., Whether such high precision spin estimation is possible in practice is almost certainly going to be limited by the systematic uncertainties in the modeling of the emission region.
 systematicNevertheless. the constraining power of even a handful of closure phase measurements is clear.," Nevertheless, the constraining power of even a handful of closure phase measurements is clear."
 The measurement of mm-VLBI closure phases in Ser A* provide an independent. significant constraint upon the structure of the emission region.," The measurement of mm-VLBI closure phases in Sgr A* provide an independent, significant constraint upon the structure of the emission region."
 Models based upon radiatively inefficient accretion flows do not generally have small closure and thus are not consistent with the recently measuredphases. value on the generallySMT-CARMA-JCMT triangle.," Models based upon radiatively inefficient accretion flows do not generally have small closure phases, and thus are not generally consistent with the recently measured value on the SMT-CARMA-JCMT triangle."
 As a even in the presence of uncertainties. closure consequence.phases can place significant constraintslarge upon the structure of Ser A*.," As a consequence, even in the presence of large uncertainties, closure phases can place significant constraints upon the structure of Sgr A*."
 However. when only accretion flow models that. fit the 1.3mm-VLBI visibility amplitudes are considered. the probability of finding an excluded closure phase falls to less than3%.," However, when only accretion flow models that fit the $1.3\mm$ -VLBI visibility amplitudes are considered, the probability of finding an excluded closure phase falls to less than."
. Hence. the vast majority of the acceptable models in Brodericketal.(2010) are broadly consistent with all current 1.3mm-VLBI constraints.," Hence, the vast majority of the acceptable models in \citet{Brod_etal:10} are broadly consistent with all current $1.3\mm$ -VLBI constraints."
 This is surprising since the 1.3mm-VLBI amplitudes and closure phase measurements provide independent constraints. and thus having fit the aceretion models to the former there is no a priori reason to expect that they would be consistent with the latter.," This is surprising since the $1.3\,\mm$ -VLBI amplitudes and closure phase measurements provide independent constraints, and thus having fit the accretion models to the former there is no a priori reason to expect that they would be consistent with the latter."
" Based upon the models that are consistent with the |.3mm-VLBI visibility amplitudes. we predict that the closure phase on the SMT-CARMA-JCMT triangle during the time it was measured was between +30"". with +13° most preferred."," Based upon the models that are consistent with the $1.3\mm$ -VLBI visibility amplitudes, we predict that the closure phase on the SMT-CARMA-JCMT triangle during the time it was measured was between $\pm30^\circ$, with $\pm13^\circ$ most preferred."
 Improving the SNR of mm-VLBI observations by a factor of a few. achievable by phasing up multiple CARMA antennas and employing the recently phased sites in Hawai. should allow detection of a non-zero closure phase on the SMT-CARMA-JCMT triangle.," Improving the SNR of mm-VLBI observations by a factor of a few, achievable by phasing up multiple CARMA antennas and employing the recently phased sites in Hawaii, should allow detection of a non-zero closure phase on the SMT-CARMA-JCMT triangle."
 Note that this would eliminate the existing 180° ambiguity in the orientation of Ser A*., Note that this would eliminate the existing $180^\circ$ ambiguity in the orientation of Sgr A*.
 If Ser A* is well described inefficient aceretion flow models. currently acceptableby radiativelyclosure phases on SMT-Hawaii-Chile and Hawaii-Chile-LMT) largecan be as triangleslarge (e.g..as 907. and typical values on triangles which include long baselines of 457.," If Sgr A* is well described by radiatively inefficient accretion flow models, currently acceptable closure phases on large triangles (e.g., SMT-Hawaii-Chile and Hawaii-Chile-LMT) can be as large as $90^\circ$, and typical values on triangles which include long baselines of $45^\circ$."
 Future improvements in sensitivity associated with the near-term development of the EHT will substantially increase the precision with which closure phases can be measured., Future improvements in sensitivity associated with the near-term development of the EHT will substantially increase the precision with which closure phases can be measured.
 On all triangles we considered (constructed from the SMT. CARMA. Hawaii. the LMT. and Chile) the resulting precision should be sufficient to measure the deviations from 07.," On all triangles we considered (constructed from the SMT, CARMA, Hawaii, the LMT, and Chile) the resulting precision should be sufficient to measure the deviations from $0^\circ$."
 Improved closure phase measurements and baseline coverage result in. dramatically improved. constraints upon the spin of Ser A*., Improved closure phase measurements and baseline coverage result in dramatically improved constraints upon the spin of Sgr A*.
 Including a single antenna in Chile (e.g.. APEX or a single ALMA dish) produces spin estimates with precisions of roughly +0.1.," Including a single antenna in Chile (e.g., APEX or a single ALMA dish) produces spin estimates with precisions of roughly $\pm0.1$."
 With a partially phased ALMA station. this improves to 00.03. and with the LMT is better than £0.01. at which point we fail to resolve the likely region.," With a partially phased ALMA station, this improves to $\pm0.03$, and with the LMT is better than $\pm0.01$, at which point we fail to resolve the likely region."
 In the latter cases. the spin estimation will almost certainly be limited by systematic uncertainties in the modeling of the emission region motivating the consideration of more sophisticated aceretion and outflow models.," In the latter cases, the spin estimation will almost certainly be limited by systematic uncertainties in the modeling of the emission region motivating the consideration of more sophisticated accretion and outflow models."
 This work was supported in part by NSF grants. AST- AST-0807843 and AST-0903844. and NASA grants NXOSALA3G and NNAO9DB30A.," This work was supported in part by NSF grants AST-0907890, AST-0807843 and AST-0905844, and NASA grants NNX08AL43G and NNA09DB30A."
As a consequence. the binary fractions are not very well determined. as reflected by the large confidence intervals for the EBT? and EHB groups.,"As a consequence, the binary fractions are not very well determined, as reflected by the large confidence intervals for the EBT2 and EHB groups."
 However. the statistical analysis still permits us to draw many interesting conelusions.," However, the statistical analysis still permits us to draw many interesting conclusions."
 The calculations of rracarebasedonthef rdc tionofde, The calculations of \\ref{c_frac} are based on the fraction of detected binaries relative to the number of stars observed in each group.
tecte «dbinariesrelativetot," Some systematic error in the estimated temperatures could thus alter the results, in particular in the presence of stars evolving off the EHB toward higher luminosities, whose temperature could be underestimated."
he pkTRUenum λα TM rereVodiagram. becausethebluersectionof theH Bisalmostvertical.," The temperature and luminosity of the hotter stars are degenerate in the $V$ vs $B-V$ ) diagram, because the bluer section of the HB is almost vertical."
 ir the «le finit. Voplane.," Small luminosity variations or temperature errors should not affect the definition of the EBT1 and EBT2 samples, because it is based on the presence of clear photometric gaps and not on temperature boundaries."
 wherethetemperature luminositvdegeneracyismuchless pronounced. andEBT 3stars. ov dirPe on," Stars ascending the asymptotic giant branch can be identified in the $V$ vs $U-V$ ) plane, where the temperature-luminosity degeneracy is much less pronounced, and EBT3 stars, evolving directly to the WD cooling sequence, should be at all stages fainter than our magnitude limit."
BAIE Fasyy Thus. uinteThan HBstarsoriginating fromthe EBT2 po f pulation.," As a consequence, targets assigned to the EBT2 sample should be correctly classified, while two EBT1 targets brighter and/or bluer than the main HB population (see he upper panel of Figure \ref{f_cmd}) ) could be evolved post-HB stars originating from the EBT2 population."
"Oneo thenweasexAsstenindirtl ber aig SOTWathun the refe,esults)) .wAiletheotheroneshowsnosigno fbinarity."," One of them was excluded from the statistical analysis for other reasons \\ref{c_results}) ), while the other one shows no sign of binarity."
 8Η] Downs [Sins referac EBT2.. causingachangeo fabout2% inthevaluesgivenin ," Assigning it to the EBT2 group, we would have $N$ =15 in \\ref{c_fracEBT2}, causing a change of about in the values given in Table \ref{t_results}."
"FADES, ry JSE"," The results for the EBT1 group are even less sensitive to the change, because of the larger sample."
RIORA help[9 ," In conclusion, the uncertainties in the definition of the EBT1 and EBT2 samples have a negligible impact on the results."
OI MenMihechange. / fona.," In contrast, even small systematics in the temperature scale can strongly affect the statistical analysis of the EHB group, defined only by a temperature limit, because of the small number of observed targets."
 And fipmay Can vary by £5% and z1566.. respectively.," For example, when $N$ is varied by $\pm$ 1, $f_\mathrm{c,max}$ and $f_\mathrm{ip,max}$ can vary by $\pm$ and $\pm$, respectively."
 These changes are not large when compared to the uncertainties given in Table 2.. and the general conclusions are unaffected. but this result warns us that the results for the EHB group should be regarded as only indicative. as already stated in refeΓΕ ΠΛ.," These changes are not large when compared to the uncertainties given in Table \ref{t_results}, and the general conclusions are unaffected, but this result warns us that the results for the EHB group should be regarded as only indicative, as already stated in \\ref{c_fracEHB}."
 Binaries with periods longer than 10 days have never been studied in GCs. and have been the target of very few surveys even among field stars?).," Binaries with periods longer than 10 days have never been studied in GCs, and have been the target of very few surveys even among field stars."
. The most probable estimate of the intermediate-period binary fraction among EHB stars is very high. but is most probably affected by small number statistics as can be deduced from the very wide probability curve in the lower panel of Figure 12..," The most probable estimate of the intermediate-period binary fraction among EHB stars is very high, but is most probably affected by small number statistics as can be deduced from the very wide probability curve in the lower panel of Figure \ref{f_frac}."
 However. an important result of our investigation is a relatively high 20%)) lower limit for fij; among EHB stars.," However, an important result of our investigation is a relatively high ) lower limit for $f_\mathrm{ip}$ among EHB stars."
 The estimated close binary fraction in all the clusters studied so far is either comparable to2808..5986.. M80)) or much lower than 6752)) this value.," The estimated close binary fraction in all the clusters studied so far is either comparable to, ) or much lower than ) this value."
" The value of fi, is most probably higher than f; even in the EBT2. because their probability curves are similar but p(fip) is shifted toward higher values."," The value of $f_\mathrm{ip}$ is most probably higher than $f_\mathrm{c}$ even in the EBT2, because their probability curves are similar but $f_\mathrm{ip}$ ) is shifted toward higher values."
 The most probable values indicate that only one-fifth of the EBT? stars can reside in close systems. but that even half of them could be binary systems with y x50 days.," The most probable values indicate that only one-fifth of the EBT2 stars can reside in close systems, but that even half of them could be binary systems with $\wp\leq$ 50 days."
 Although we cannot provide a more reliable estimate. this result suggests that binaries wider than those investigated so far could play an important role and should deserve more attention. particularly where close systems are lacking.," Although we cannot provide a more reliable estimate, this result suggests that binaries wider than those investigated so far could play an important role and should deserve more attention, particularly where close systems are lacking."
 A second prominent result clearly evident from Table 2 is that the confidence intervals in the EBT! and EBT2 do not overlap. for both f. and fi.," A second prominent result clearly evident from Table \ref{t_results} is that the confidence intervals in the EBT1 and EBT2 do not overlap, for both $f_\mathrm{c}$ and $f_\mathrm{ip}$."
 This means that the probability that the binary fraction has the same value 1 these two sections of the HB ts negligible (of the order of 190):fraction. there being a very small quantity of binaries among the stars cooler than the gap. and about for hotter objects.," This means that the probability that the binary fraction has the same value in these two sections of the HB is negligible (of the order of ):, there being a very small quantity of binaries among the stars cooler than the gap, and about for hotter objects."
 Unfortunately. our data cannot exclude f. and fiy monotonically increasing with the temperature rather than WdDS ys uusaL erg," Unfortunately, our data cannot exclude $f_\mathrm{c}$ and $f_\mathrm{ip}$ monotonically increasing with the temperature rather than showing a real discontinuity at G1."
 nh. hhnt Ibis. thesities NrgM WPheestimate. ΙΟ. Unt ο ow pyA EBT? and in its (hotter) sub-group EHB. in particular for tae for Din doe eheyshaegsu aha hs dAUS ers fnb lo ," In this regard, we note that there is reasonable agreement between the results in the EBT2 and in its (hotter) sub-group the EHB, in particular for the close binary fraction (compare also the shapes of the dotted curves in the middle and lower panels of Figure \ref{f_frac}) )."
"ADR, MOSS a, UI a(un poc Wu EBT? population. ‘instead of increasiο with temperature. but no firm conclusion INT. = fet: WA ieAGEn"," Thus, $f_\mathrm{c}$ could be quite homogeneous within the EBT2 population, instead of increasing with temperature, but no firm conclusion can be drawn because of the too wide confidence intervals."
 Future observations should us to clarify this issue: a sudden increase in the binarity in correspondence with Gl would strongly relate it to the formation of all EBT2 stars. while a f. slowly 1creasing with temperature may indicate that the progeny of close systems are preferentially hotter.," Future observations should help us to clarify this issue: a sudden increase in the binarity in correspondence with G1 would strongly relate it to the formation of all EBT2 stars, while a $f_\mathrm{c}$ slowly increasing with temperature may indicate that the progeny of close systems are preferentially hotter."
 In the context of the He-enrichment scenario for. the formatior of hot HB stars in GCs. the difference in. binary fraction among EBT] and EBT? stars contrasts with the investigation of inM4.. who found a much lower quantity of binaries among red giants that displayed evidence of chemical enrichment. than among normal ones.," In the context of the He-enrichment scenario for the formation of hot HB stars in GCs, the difference in binary fraction among EBT1 and EBT2 stars contrasts with the investigation of in, who found a much lower quantity of binaries among red giants that displayed evidence of chemical enrichment, than among normal ones."
 They argue that this difference is naturally explained by assuming that the second stellar generation formed in a denser environment. where more frequent dynamical interactions enhanced the disruption rate of binary systems.," They argue that this difference is naturally explained by assuming that the second stellar generation formed in a denser environment, where more frequent dynamical interactions enhanced the disruption rate of binary systems."
 However. in the multi-population models of2808.. the EBT2 is interpreted as the progeny of the latest and He-richest of the three populations observed in the MS and. following?.. we would expect the EBT2 to be depleted in binaries. at variance with what is observed.," However, in the multi-population models of, the EBT2 is interpreted as the progeny of the latest and He-richest of the three populations observed in the MS and, following, we would expect the EBT2 to be depleted in binaries, at variance with what is observed."
 If their results were confirmed as a general behavior of the chemically polluted stars in GCs. our results would argue against the link between the EBT2 and the He-enriched stars.," If their results were confirmed as a general behavior of the chemically polluted stars in GCs, our results would argue against the link between the EBT2 and the He-enriched stars."
 However. an alternative interpretation would be that both the He-enrichment and the binary scenarios co-exist in the cluster. as different channels for the formation of blue HB stars.," However, an alternative interpretation would be that both the He-enrichment and the binary scenarios co-exist in the cluster, as different channels for the formation of blue HB stars."
 In this case. in the EBT? both the progeny of He-enriched stars and products of binary interactions would be found.," In this case, in the EBT2 both the progeny of He-enriched stars and products of binary interactions would be found."
 This would cause a higher frequency of EBT2 stars in the HB with respect to the fraction of He-rich MS stars. but binaries in GCs usually represent aminor fraction of the entire population. and this difference could pass unnoticed.," This would cause a higher frequency of EBT2 stars in the HB with respect to the fraction of He-rich MS stars, but binaries in GCs usually represent aminor fraction of the entire population, and this difference could pass unnoticed."
 No splitting of the red- or sub- giant branch has been detected so far in 2808.. but this cluster shows a strong Na-O anticorrelation (?).. which has often been," No splitting of the red- or sub- giant branch has been detected so far in , but this cluster shows a strong Na-O anticorrelation , which has often been"
the results from 2011 January 22.,the results from 2011 January 22.
 We do not detect any variability among any of the epochs., We do not detect any variability among any of the epochs.
" Aside from multiplicative scaling (+10%)) representative of typical calibration uncertainties, the spectra from the three epochs have no significant differences."," Aside from multiplicative scaling $\pm 10$ ) representative of typical calibration uncertainties, the spectra from the three epochs have no significant differences."
" Although the absolute calibration is uncertain at a ~10% level, the relative magnitude of the amplitude parameters listed in Table 2 are precise at a level of or better."," Although the absolute calibration is uncertain at a $\approx$ level, the relative magnitude of the amplitude parameters listed in Table \ref{table} are precise at a level of or better."
 The magnitude and precision of the fitted line centers and widths are unaffected by amplitude scaling., The magnitude and precision of the fitted line centers and widths are unaffected by amplitude scaling.
" The ammonia spectral features at vpsg5—50, —49, and —4T7kms-! represent the first maser species in NGC 7538, besides water, to be observed having vrsn> —51kms-!."," The ammonia spectral features at $v_{\rm LSR} \approx -50$, $-49$, and $-47\,{\rm km}\,{\rm s}^{-1}$ represent the first maser species in NGC 7538, besides water, to be observed having $v_{\rm LSR} > -51\,{\rm km}\,{\rm s}^{-1}$ ."
 These weak features are not necessarily associated with IRS 1., These weak features are not necessarily associated with IRS 1.
" At 18 GHz, the GBT has a —33.0- sidelobe at the location of IRS 9 when pointed at IRS 1."," At 18 GHz, the GBT has a $-33.0$ -dB sidelobe at the location of IRS 9 when pointed at IRS 1."
" For example, the 20-mJy signal in Figure 2 may be a 40-Jy ammonia maser in IRS 9."," For example, the 20-mJy signal in Figure \ref{gbt} may be a 40-Jy ammonia maser in IRS 9."
" A large amplitude for an ammonia maser is not without precedent: Madden et ((1986) and Pratap et ((1991) find the (9,6) maser in W51 to have a flux density greater than 50 Jy."," A large amplitude for an ammonia maser is not without precedent: Madden et (1986) and Pratap et (1991) find the (9,6) maser in W51 to have a flux density greater than 50 Jy."
 IRS 11 lies in a null of the 18-GHz beam pattern and likely does not contribute to the observed signal at 18.5 GHz., IRS 11 lies in a null of the 18-GHz beam pattern and likely does not contribute to the observed signal at 18.5 GHz.
 We are planning EVLA observations in order to precisely image the locations of all of the ammonia spectral features., We are planning EVLA observations in order to precisely image the locations of all of the ammonia spectral features.
(3-2) and (10-9).,(3-2) and (10-9).
 In these tracers we detect overall 8 cores., In these tracers we detect overall 8 cores.
" The coordinates and the sizes ol the cores are given in Table 2) and the column densities of the observed species towards (he cores are given in Table ο,", The coordinates and the sizes of the cores are given in Table \ref{coord} and the column densities of the observed species towards the cores are given in Table \ref{column}.
 The cores are better defined in the higher spatial resolution maps. ie. (1-0) and (10-9).," The cores are better defined in the higher spatial resolution maps, i.e. (1-0) and (10-9)."
 Cores Lupl C1 and Lupl C2 are visible in the CS and maps. only mareinally in NIL;. but not in the other species.," Cores Lup1 C1 and Lup1 C2 are visible in the CS and maps, only marginally in $_3$, but not in the other species."
 Core Lupl C7 is not detected in (3-2) and only marginally in the higher (10-9) line., Core Lup1 C7 is not detected in (3-2) and only marginally in the higher (10-9) line.
 On the contrary. Lupl C8 is detected in and CS but not in NI; and.," On the contrary, Lup1 C8 is detected in and CS but not in $_3$ and."
. Similarly. Lupl C5 is detected in (3-2) but nol in δα it has not been mapped in.," Similarly, Lup1 C5 is detected in (3-2) but not in $_3$; it has not been mapped in."
. The ΕΝ (2-2) main component in Lupl C8 show an asvnuuetric line profile (see Fig. 5)), The $_3$ N (3-2) main component in Lup1 C8 show an asymmetric line profile (see Fig. \ref{hc3n_c8}) )
 that could be a sign of infall., that could be a sign of infall.
 ILowever. ad," However, at"
elements are more efficiently retained in the cluster than the faster moving Type II ejecta (Smith et al. 2000)).,elements are more efficiently retained in the cluster than the faster moving Type II ejecta (Smith et al. \cite{smith00}) ).
" These authors also suggest that all elements produced by Type Ia SNe escape from the ω CCen cluster. which would explain the ""high"" trend of ffound by Johnson Pilachowski (2010))."," These authors also suggest that all elements produced by Type Ia SNe escape from the $\omega$ Cen cluster, which would explain the `high' trend of found by Johnson Pilachowski \cite{johnson10}) )."
 As seen from this discussion. only aand agree between the stars in the co CCen cluster and those in the low-a population.," As seen from this discussion, only and agree between the stars in the $\omega$ Cen cluster and those in the $\alpha$ population."
 There are large differences in the distributions of[a/Fe].[Na/Fe]. and [Ba/Y].," There are large differences in the distributions of, and ."
 Hence. if some of the low-c stars were accreted from ω CCen. the chemical evolution pattern must have been different in the inner and outer parts of the progenitor galaxy.," Hence, if some of the $\alpha$ stars were accreted from $\omega$ Cen, the chemical evolution pattern must have been different in the inner and outer parts of the progenitor galaxy."
 Evidence of such à scenario comes from the recent work by Chou et al. (2010)), Evidence of such a scenario comes from the recent work by Chou et al. \cite{chou10}) )
 on abundances in the Sagittarius dSph galaxy: stars in the leading north arm tend to have lower values of [La/Y] than stars in the Ser core for the metallicity range —1<[Fe/H]0 (see Fig., on abundances in the Sagittarius dSph galaxy; stars in the leading north arm tend to have lower values of [La/Y] than stars in the Sgr core for the metallicity range $-1 < \feh < 0$ (see Fig.
 6 in the Chou et al., 6 in the Chou et al.
 paper)., paper).
 In this paper very precise abundances of Mn. Cu. Zn. Y. and Ba have been determined relative to Fe for stars in the solar neighborhood with halo kinematics.," In this paper very precise abundances of Mn, Cu, Zn, Y, and Ba have been determined relative to Fe for stars in the solar neighborhood with halo kinematics."
 The two populations. high- and low-a stars. that were found to have different ratios of[a/Fe].[Na/Fe]. and iin our first paper (NSIO) are also separated in[Cu/Fe].[Zn/Fe]. and |Ba/Y].," The two populations, high- and $\alpha$ stars, that were found to have different ratios of, and in our first paper (NS10) are also separated in, and ."
 There is. however. no significant differences inMn/Fe|.," There is, however, no significant differences in."
 Trends and correlations between the abundance ratios may to a large extent be explained from existing nucleosynthesis calculations if the high-a stars have been formed in regions with such a high SER that only massive stars and IID have contributed to the chemical evolution up to [Fe/H]=—0.4., Trends and correlations between the abundance ratios may to a large extent be explained from existing nucleosynthesis calculations if the $\alpha$ stars have been formed in regions with such a high SFR that only massive stars and II have contributed to the chemical evolution up to $\feh \simeq -0.4$.
 The low-a stars. on the other hand. may originate from systems with a slower chemical evolution. characterized by delayed enrichment from relatively metal-poor IIa and low-mass AGB stars in addition to the contributions from massive stars and Type II SNe.," The $\alpha$ stars, on the other hand, may originate from systems with a slower chemical evolution, characterized by delayed enrichment from relatively metal-poor Ia and low-mass AGB stars in addition to the contributions from massive stars and Type II SNe."
 It is. however. difficult to explain that there is no significant difference in bbetween the high- and low- populations.," It is, however, difficult to explain that there is no significant difference in between the high- and $\alpha$ populations."
 Furthermore. the data for ccall for a revision of the Ni yields from IIa. RGB stars in present-day dSph satellite galaxies show abundance trends of [o/Fe].[Na/Fe].[Ni/Fe]. |Cu/Fe]. and tthat have some similarities with the trends for the low-c population. but there is no exact matching.," Furthermore, the data for call for a revision of the Ni yields from Ia. RGB stars in present-day dSph satellite galaxies show abundance trends of , and that have some similarities with the trends for the $\alpha$ population, but there is no exact matching."
 In general. existing dSph galaxies seem to be characterized by an even slower chemical evolution than the systems in which the low-a stars were formed.," In general, existing dSph galaxies seem to be characterized by an even slower chemical evolution than the systems in which the $\alpha$ stars were formed."
 As discussed in NS10. the kinematics of the low- stars suggest that some of them were acereted from the progenitor galaxy of the ω CCen globular cluster.," As discussed in NS10, the kinematics of the $\alpha$ stars suggest that some of them were accreted from the progenitor galaxy of the $\omega$ Cen globular cluster."
 We have looked for supporting evidence from chemical abundance ratios. but have found more differences between the w CCen cluster and the low-« stars than similarities.," We have looked for supporting evidence from chemical abundance ratios, but have found more differences between the $\omega$ Cen cluster and the $\alpha$ stars than similarities."
 The distributions of aand ffor RGB stars in «o CCen overlap well with those of the low-w population. but aand aare very different in ω CCen and low-a stars.," The distributions of and for RGB stars in $\omega$ Cen overlap well with those of the $\alpha$ population, but and are very different in $\omega$ Cen and $\alpha$ stars."
 This difference could be explained if the products of AGB stars are selectively retained within the ω CCen cluster., This difference could be explained if the products of AGB stars are selectively retained within the $\omega$ Cen cluster.
 The trend of iis also different for «o CCen and the low-c population. so one has to invoke differential loss of IIL and Ha winds from ω CCen if some of the low-a stars did originate from the progenitor galaxy.," The trend of is also different for $\omega$ Cen and the $\alpha$ population, so one has to invoke differential loss of II and Ia winds from $\omega$ Cen if some of the $\alpha$ stars did originate from the progenitor galaxy."
 Zolotov et al. (2009.. 2010))," Zolotov et al. \cite{zolotov09}, \cite{zolotov10}) )"
 have recently used N-body and smooth particle hydrodynamic simulations to investigate the kinematics and ttrends of stellar halos of large galaxies similar to the Milky Way., have recently used N-body and smooth particle hydrodynamic simulations to investigate the kinematics and trends of stellar halos of large galaxies similar to the Milky Way.
 They find that the inner halos (R<20 kkpe) contain both accreted and in situ formed stars., They find that the inner halos $R < 20$ kpc) contain both accreted and in situ formed stars.
 These two populations are separated in aat the high end of the metallicity distribution function if the halo was formed in connection with a few major mergers CMattic/Mpiinay> 0.1) at early times. i.e. more than 8 — 9 Gyr ago.," These two populations are separated in at the high end of the metallicity distribution function if the halo was formed in connection with a few major mergers $M_{\rm satellite}/ M_{\rm primary} \, > 0.1$ ) at early times, i.e. more than 8 – 9 Gyr ago."
 The in situ halo stars formed in the innermost ~4 kkpc of the galaxy in a deep potential well causing the SFR to be so high that only core collapse SNe contributed to the chemical enrichment., The in situ halo stars formed in the innermost $\sim 4$ kpc of the galaxy in a deep potential well causing the SFR to be so high that only core collapse SNe contributed to the chemical enrichment.
" Later these stars were ""heated"" to halo kinematics by mergers.", Later these stars were `heated' to halo kinematics by mergers.
 The accreted stars. on the other hand. formed in satellite galaxies with shallower potential wells and hence lower SER allowing Ila to contribute to the enrichment at the high end of the metallicity distribution.," The accreted stars, on the other hand, formed in satellite galaxies with shallower potential wells and hence lower SFR allowing Ia to contribute to the enrichment at the high end of the metallicity distribution."
 The dual distribution of oobtained in the simulations of Zolotov et al. (2010)), The dual distribution of obtained in the simulations of Zolotov et al. \cite{zolotov10}) )
 corresponds qualitatively to the abundance trends determined in NS10 and in this paper., corresponds qualitatively to the abundance trends determined in NS10 and in this paper.
 Thus. the high-c population may consist of stars formed in the innermost part of the MilkyWay and displaced to the halo by mergers. whereas the low-« stars may have been accreted at early times from a few. relatively massive satellite," Thus, the $\alpha$ population may consist of stars formed in the innermost part of the MilkyWay and displaced to the halo by mergers, whereas the $\alpha$ stars may have been accreted at early times from a few, relatively massive satellite"
develop once the galaxy halos begin to interpenctrate at close separation.,develop once the galaxy halos begin to interpenetrate at close separation.
 While ai detailed. analysis of the dilluse gaseous intragroup medium is bevond the primary focus of our work. the basic properties of this component appear to be consistent with observations.," While a detailed analysis of the diffuse gaseous intragroup medium is beyond the primary focus of our work, the basic properties of this component appear to be consistent with observations."
" In. particular. our moclel predicts that. at the present state. the intragroup medium. is predominantly warn. LO’10 Ex. has a fairly: low density. 10!"" and fills the entire volume that we simulate."," In particular, our model predicts that, at the present state, the intragroup medium is predominantly warm $10^5-10^6$ K, has a fairly low density, $10^{-4} -
10^{-6}$ $^{-3}$, and fills the entire volume that we simulate."
" We note this medium is also expected. to extend to larger scales into what has been termed the ""Marm-LHot Intergalactic Medium"" (?77).."," We note this medium is also expected to extend to larger scales into what has been termed the “Warm-Hot Intergalactic Medium"" \citep{CO99,HGM98,Dav01}."
 While the presence. of our intragroup medium is consistent with all current data (?).., While the presence of our intragroup medium is consistent with all current data \citep{Oso02}.
 Owing to the dillieulty in directly observing gas in this state. most evidence for its existence comes from indirect. means.," Owing to the difficulty in directly observing gas in this state, most evidence for its existence comes from indirect means."
 For example. Chandra and £USL observations of 2.~0 Oxveen anc Neon absorption along numerous sight. lines sugeests the presence ofa local. volumefilling. diffuse. warn medium (2277).. although detailed analysis of the ionization states suggest that this is a complex multiphase medium that our simulation does not have the resolution to model.," For example, $Chandra$ and $FUSE$ observations of $z\sim0$ Oxygen and Neon absorption along numerous sight lines suggests the presence of a local, volume–filling, diffuse, warm medium \citep{Nic02,Nic03,Sem03,Sav03}, although detailed analysis of the ionization states suggest that this is a complex multi–phase medium that our simulation does not have the resolution to model."
 While Figure 2) presented the dynamical evolution of the stellar mass on large scales. this vantage point makes it cillicult to distinguish the tell-tale signs of a galaxy merger.," While Figure \ref{fig:starimages} presented the dynamical evolution of the stellar mass on large scales, this vantage point makes it difficult to distinguish the tell-tale signs of a galaxy merger."
 We therefore zoom into the central regions of the Local Group ancl specifically show the merger between the Milky Wav and Andromeda in Figure ον, We therefore zoom into the central regions of the Local Group and specifically show the merger between the Milky Way and Andromeda in Figure \ref{fig:zoomimage}.
 From this viewpoint the classic signatures of a galaxy interaction. such as tidal tails. plumes. and shells are clearly. evident.," From this viewpoint the classic signatures of a galaxy interaction, such as tidal tails, plumes, and shells are clearly evident."
 The physical separation between the Milkv. Way and Andromeda is presented. in. Figure 5.., The physical separation between the Milky Way and Andromeda is presented in Figure \ref{fig:csep}.
 “Phe current. state of the Local Group occurs 5 Gyr after the start of our, The current state of the Local Group occurs $\sim5$ Gyr after the start of our
The radio emission can be attributed to a jet on accoun of its self-absorbed. spectrum. (Fender. 2006 ancl references herein).,The radio emission can be attributed to a jet on account of its self-absorbed spectrum (Fender 2006 and references therein).
 I£ the jet were to dominate the optical and infrarec emission then we would expect to see a greater correlation »etween their respective lighteurves., If the jet were to dominate the optical and infrared emission then we would expect to see a greater correlation between their respective lightcurves.
 Indeed. the lighteurves shown here could. lead us to assume that the optical anc infared are dominated by disc emission were it not lor he additional svnchrotron components found by Zurita ο al. (," Indeed, the lightcurves shown here could lead us to assume that the optical and infared are dominated by disc emission were it not for the additional synchrotron components found by Zurita et al. ("
2006). LIvnes et al. (,"2006), Hynes et al. ("
2006) ancl Alaitra ct al. (,2006) and Maitra et al. (
2009).,2009).
 eonetheless.. despite the known additional. synchrotron component. the excess of optical/infrared emission following he X-rav/radio decay shows that the cise emission is still ikelv to be significant at these wavelengths.," Nonetheless, despite the known additional synchrotron component, the excess of optical/infrared emission following the X-ray/radio decay shows that the disc emission is still likely to be significant at these wavelengths."
 There has been recent debate regarding whether there can be a disc component to the X-ray. spectrum when a source is in the lowhard state (e.g. Reis ct al., There has been recent debate regarding whether there can be a disc component to the X-ray spectrum when a source is in the low/hard state (e.g. Reis et al.
 2009: Rvkoll et al., 2009; Rykoff et al.
 2007: Miller et al., 2007; Miller et al.
 2006: Chiang et al., 2006; Chiang et al.
 2009)., 2009).
 Cüerli&sski. Done Page (2008. 2009) have explained the softening of the spectrum in terms of irradiation of the inner disc and Done Díaaz Trigo (2009) demonstrated. that the iron line is an artefact of pile-up. thereby. resolving the controversy over whether or not the disc is truncated.," Gierlińsski, Done Page (2008, 2009) have explained the softening of the spectrum in terms of irradiation of the inner disc and Done Díaaz Trigo (2009) demonstrated that the iron line is an artefact of pile-up, thereby resolving the controversy over whether or not the disc is truncated."
 This is also discussed in Alaitra et al. (, This is also discussed in Maitra et al. (
2009). whose mocelling of one epoch during the decay of NTE J11IS|480 in. 2005 allows for (although does not require) a small inner disc radius and. consequently. a source of soft N-ravs with which to irradiate the outer disc.,"2009), whose modelling of one epoch during the decay of XTE J1118+480 in 2005 allows for (although does not require) a small inner disc radius and, consequently, a source of soft X-rays with which to irradiate the outer disc."
 For à typical low/harcl state outburst. we would expect the power-law component of the X-ray emission to dominate at both high and low energies and for there to be little difference between the respective lighteurves., For a typical low/hard state outburst we would expect the power-law component of the X-ray emission to dominate at both high and low energies and for there to be little difference between the respective lightcurves.
 That indeed appears to be the case in the high and low energy. N-ray lighteurves presented here and there is no suggestion that there are contributions [rom more than one component (contrasting with the liehtcurves of. for example. GRO 40: Brocksopp ct al.," That indeed appears to be the case in the high and low energy X-ray lightcurves presented here and there is no suggestion that there are contributions from more than one component (contrasting with the lightcurves of, for example, GRO $-$ 40; Brocksopp et al."
 2006)., 2006).
 Once the X-ravs and radio have started to decay. their behaviour is closely linked.," Once the X-rays and radio have started to decay, their behaviour is closely linked."
" Γον all appear to undergo some aclelitional event around ALJD 5338553390. manifested as a ""shoulder"" superimposed on the decay."," They all appear to undergo some additional event around MJD 53385–53390, manifested as a “shoulder” superimposed on the decay."
 The profile of the radio lightcurve is more comparable with that of sources which made transitions to softer spectral states. such as ATE 318 or 00 (Drocksopp et al.," The profile of the radio lightcurve is more comparable with that of sources which made transitions to softer spectral states, such as XTE $-$ 318 or $-$ 00 (Brocksopp et al.,"
 2005: κους et al., 2005; Kuulkers et al.
 1999. respectively) than we would. expect for a source in the low/hared state., 1999 respectively) than we would expect for a source in the low/hard state.
" Ao second. optically-thin. jet ejection event seems a likely cause of this radio ""shoulder. particularly as the radio spectrum at this time is particularly poorly fit hy a power-law."," A second, optically-thin, jet ejection event seems a likely cause of this radio “shoulder”, particularly as the radio spectrum at this time is particularly poorly fit by a power-law."
 Confirmation of this would. require detection of a simultaneous reduction in the spectral index. needing higher sensitivity. and. sampling in the radio data.," Confirmation of this would require detection of a simultaneous reduction in the spectral index, needing higher sensitivity and sampling in the radio data."
 Multiple ejections are commonly. associated with transient events. (Drocksopp et al., Multiple ejections are commonly associated with transient events (Brocksopp et al.
 2002) but more usually with the optically thin events of sources which soften and enter the very high. state (Fender ct al., 2002) but more usually with the optically thin events of sources which soften and enter the very high state (Fender et al.
 2004)., 2004).
 Such jet ejections may be unexpected. during the lowπαντα state but not unprecedented (e.g. GS 64 DBrocksopp et al., Such jet ejections may be unexpected during the low/hard state but not unprecedented (e.g. GS $-$ 64 Brocksopp et al.
 2001)., 2001).
 It may be that the “canonical” stable. compact Lat-spectrum jet of the low/hard state (e.g. (νο N-1. GN 339. 4 or the 2000 outburst of NPE J111IS|480: Markolf al.," It may be that the “canonical” stable, compact flat-spectrum jet of the low/hard state (e.g. Cyg X-1, GX $-$ 4 or the 2000 outburst of XTE J1118+480; Markoff et al."
 2005 ancl references therein) may turn out to be the etexception rather than the norm., 2005 and references therein) may turn out to be the exception rather than the norm.
 Alternatively. this 7ejection mia be more like the [are seen in V404 Cve during quiescence. thought to be some sort. of re-energising of the electrons within the jet rather than a new ejection event CMiller-Jones et al.," Alternatively, this “ejection” may be more like the flare seen in V404 Cyg during quiescence, thought to be some sort of re-energising of the electrons within the jet rather than a new ejection event (Miller-Jones et al."
 2008). albeit on a much longer timescale.," 2008), albeit on a much longer timescale."
 Given the link between the power-law X-ray. emission and the jet it might seem surprising that the rellare at —NLID 53415 was detected: at optical ancl hard X-ray wavelengths but not the radio., Given the link between the power-law X-ray emission and the jet it might seem surprising that the reflare at $\sim$ MJD 53415 was detected at optical and hard X-ray wavelengths but not the radio.
 We note that the X-ray reHlare preceded: the optical rellare and. occurred in a gap between radio observations and so we could have missed a similar event in the radio due to the sparse sampling., We note that the X-ray reflare preceded the optical reflare and occurred in a gap between radio observations and so we could have missed a similar event in the radio due to the sparse sampling.
 Alultiwavelength daily monitoring of these events is required at high sensitivity to determine their true nature., Multiwavelength daily monitoring of these events is required at high sensitivity to determine their true nature.
 Finally. we compare these results with the X-ray. optical and radio lighteurves of the 2000 outburst of NTE ΤΙΝ|480. detailed: analyses of which were presented: by Chaty ct al. (," Finally, we compare these results with the X-ray, optical and radio lightcurves of the 2000 outburst of XTE J1118+480, detailed analyses of which were presented by Chaty et al. ("
2003) and Brocksopp et al. (,2003) and Brocksopp et al. (
2004).,2004).
 The lighteurves. of that event are notable for their highly correlated behaviour at all frequencies. with the long.," The lightcurves of that event are notable for their highly correlated behaviour at all frequencies, with the long,"
"Alassive. stars dominate. the uu""light. and. kinematics. of ⋅⊲⊀⋅⊲⋎voung and star-forming⋅. galaxies.",Massive stars dominate the light and kinematics of young and star-forming galaxies.
" They spend most of their lifetimes with: elfective⋅⋠ temperatures greater than 107 lieWh. and hence most of their MNlight is. emitted. in. the ultraviolet. (UV)"" region. of. the clectromagnetic. spectrum."," They spend most of their lifetimes with effective temperatures greater than $10^4$ K, and hence most of their light is emitted in the ultraviolet (UV) region of the electromagnetic spectrum."
 It is ⊀⋠⋅dillicult to observe emission in the [ασV. (1000-20 AA)) from sources in the local Urdiverse. due to absorption in the Earth’s atmosphere and the limitations of space-based instruments sensitive to this waveeneth range.," It is difficult to observe emission in the far-UV $\sim$ ) from sources in the local Universe, due to absorption in the Earth's atmosphere and the limitations of space-based instruments sensitive to this wavelength range."
 At redshifts above two this emission moves into the range of optical spectroscopy., At redshifts above two this emission moves into the range of optical spectroscopy.
 “The study of rest-frame UV. spectral lines therefore become relatively straigitforward ancl provides a direct insight into the massive scllar population of forming and ultraviolet-Iuminous galaxies., The study of rest-frame UV spectral lines therefore become relatively straightforward and provides a direct insight into the massive stellar population of star-forming and ultraviolet-luminous galaxies.
 Llowever a problem: arises from the extreme distance ancl consequent apparent faintness of galaxies at such redshifts., However a problem arises from the extreme distance and consequent apparent faintness of galaxies at such redshifts.
 “Phe low signal-to-noise obtained in observations ob⋅ typical⊀ individual... galaxies. makes their aeline strength and stellar population. properties. dillieult⊀⋅ to determine⊀ with. any degree. of⋅ precision., The low signal-to-noise obtained in observations of typical individual galaxies makes their line strength and stellar population properties difficult to determine with any degree of precision.
a One approach to. circumventing: this: problem ancl characterisingMN a galaxy population. is. to combine a number of⋅ spectra to obtain⊀ a single.. mean composite spectrum.," One approach to circumventing this problem and characterising a galaxy population is to combine a number of spectra to obtain a single, mean composite spectrum."
" Studies of such stacked. spectra. notably the composite constructec by Shapleyctal.(2003) of z~3 rest-UV selected ""Lyman-break galaxies! (LBCs). allow the determination of absorption and emission line intensities that are not detectable in. individual sources. while being limited to vielding information on a ‘typical’ source rather than anv individual galaxy."," Studies of such stacked spectra, notably the composite constructed by \citet{shapley} of $z\sim3$ rest-UV selected `Lyman-break galaxies' (LBGs), allow the determination of absorption and emission line intensities that are not detectable in individual sources, while being limited to yielding information on a `typical' source rather than any individual galaxy."
 However. the analysis of Shapleyetal.(2003). revealed problems with our understanding of the stellar populations in such svstems.," However, the analysis of \citet{shapley} revealed problems with our understanding of the stellar populations in such systems."
 The standard. stellar. population models produced by. (Leithereretal.L999) were not able to simultaneously, The standard stellar population models produced by \citep{1999ApJS..123....3L} were not able to simultaneously
»* seen from Fie3 that the vast. majority (0454) of the SDSS quasars are acercting material at. rates. of AL.<M10M. per vear.,be seen from Fig\ref{fig3} that the vast majority ) of the SDSS quasars are accreting material at rates of $\Msolar~<~\Mdot~<~10\Msolar$ per year.
" However. perhaps the most striking feature of Fig 3 is the near total lack of quasars with black-hole masses of Ady),10M. ancl accretion rates close to the Eedcington limit."," However, perhaps the most striking feature of Fig 3 is the near total lack of quasars with black-hole masses of $M_{bh}>10^{9}\Msun$ and accretion rates close to the Eddington limit."
 TFhis feature. must. be regarded as significant because it is extremely unlikely tha large numbers of such objects are missing from the SDSS quasar sample., This feature must be regarded as significant because it is extremely unlikely that large numbers of such objects are missing from the SDSS quasar sample.
 We note here that this situation is at. leas qualitatively consistent with models of black-hole &rowth in which the exponential. Exldington-limited. growth of the Xdack-hole is terminated by a physical limit to the amount of material which can be supplied for accretion (eg.," We note here that this situation is at least qualitatively consistent with models of black-hole growth in which the exponential, Eddington-limited, growth of the black-hole is terminated by a physical limit to the amount of material which can be supplied for accretion (eg."
 Archibale et al., Archibald et al.
 2002: Granato ct al., 2002; Granato et al.
 2003. 2001).," 2003, 2001)."
 In. models. of this vpe. following the end of the exponential-growth phase the Mack hole is free to remain active as a luminous quasar for ~10° vears. accreting below the Eddington limit. without eading to the production of large numbers of black holes with masses 7LOMAL. which are not observed locally.," In models of this type, following the end of the exponential-growth phase the black hole is free to remain active as a luminous quasar for $\simeq 10^{8}$ years, accreting below the Eddington limit, without leading to the production of large numbers of black holes with masses $>10^{10}\Msun$ which are not observed locally."
 This ype of scenario is also consistent with the results shown in Fig 1 Fig 2. and is cliscussecd in more detail in Section 5.," This type of scenario is also consistent with the results shown in Fig 1 Fig 2, and is discussed in more detail in Section 5."
 Following the ciscovery that supermassive black holes appear to be ubiquitous at the centres of massive local ealaxies. it has been of interest to use this information to estimate the mass function of dormant. black holes in the local Universe.," Following the discovery that supermassive black holes appear to be ubiquitous at the centres of massive local galaxies, it has been of interest to use this information to estimate the mass function of dormant black holes in the local Universe."
 From a theoretical perspective it was shown bv Soltan (1982) that the local mass density of dormant slack holes could. be caleulated from the total raciated energy of optical quasars. which in turn can be calculated rom the quasar source-count distribution.," From a theoretical perspective it was shown by Soltan (1982) that the local mass density of dormant black holes could be calculated from the total radiated energy of optical quasars, which in turn can be calculated from the quasar source-count distribution."
 Traclitionally. such estimates predict. local black-hole mass densities a actor of a [few lower (Soltan 1982: Chokshi Turner 1992: Yu ‘Tremaine 2002) than those estimated. from the X-rav background (eg., Traditionally such estimates predict local black-hole mass densities a factor of a few lower (Soltan 1982; Chokshi Turner 1992; Yu Tremaine 2002) than those estimated from the X-ray background (eg.
 Fabian Lwasawa 1999). suggesting hat the dominant fraction of the quasar population is optically obscured.," Fabian Iwasawa 1999), suggesting that the dominant fraction of the quasar population is optically obscured."
 Given that a principal objective of quasar evolution models is to simultaneously. explain the evolution of both the quasar luminosity and black-hole mass functions. it is clearly important to have a robust measurement of the form of the local mass function of dormant black holes.," Given that a principal objective of quasar evolution models is to simultaneously explain the evolution of both the quasar luminosity and black-hole mass functions, it is clearly important to have a robust measurement of the form of the local mass function of dormant black holes."
" Previous estimates of the local black-hole mass function in the literature. have relied. on either the AM,{οι correlation to make the transformation [rom the local ealaxy luminosity function (e.g. Salucci et al.", Previous estimates of the local black-hole mass function in the literature have relied on either the $M_{bh}-L_{bulge}$ correlation to make the transformation from the local galaxy luminosity function (e.g. Salucci et al.
 1999). or the correlation between radio luminosity ancl black-hole mass observed in local earlv-type galaxies (c.g. Franceschini et al.," 1999), or the correlation between radio luminosity and black-hole mass observed in local early-type galaxies (e.g. Franceschini et al."
 L998)., 1998).
" More recently. both Yu Tremaine (2002) ancl Aller Richstone (2002) exploited the tight Mi,σ correlation to derive an estimate of the local mass density of dormant black holes."," More recently, both Yu Tremaine (2002) and Aller Richstone (2002) exploited the tight $M_{bh}-\sigma$ correlation to derive an estimate of the local mass density of dormant black holes."
 While the Yu Tremaine (2002) calculation was based directly on the SDSS stellar-velocity dispersion function. the Aller Richstone (2002) calculation used the Faber-Jackson relation (Faber Jackson 1976) o first. estimate the dispersion function from the galaxy uminositv function.," While the Yu Tremaine (2002) calculation was based directly on the SDSS stellar-velocity dispersion function, the Aller Richstone (2002) calculation used the Faber-Jackson relation (Faber Jackson 1976) to first estimate the dispersion function from the galaxy luminosity function."
 Despite this difference in. approach. roth studies found. consistent results for the local black-role mass density (early|late. types) with estimates of (2.120.8)-M. 10Alpe and (2940.5).10M. Alpe* or Aller Richstone and Yu ‘Tremaine respectively (444=70 km ‘Alpe ty.," Despite this difference in approach, both studies found consistent results for the local black-hole mass density (early+late types) with estimates of $(2.4\pm 0.8) \times 10^{5}\Msun$ $^{-3}$ and $(2.9\pm 0.5) \times 10^{5}\Msun$ $^{-3}$ for Aller Richstone and Yu Tremaine respectively $H_{0}=70$ km $^{-1}$ $^{-1}$ )."
 In Section 5 we will proceed to estimate the activation action of supermassive black holes ato.=2., In Section 5 we will proceed to estimate the activation fraction of supermassive black holes at $z\simeq 2$.
 This calculation requires a knowledge of the actual functional orm of the local dormant black-hole mass function. information which is unavailable from the previous studies in he literature.," This calculation requires a knowledge of the actual functional form of the local dormant black-hole mass function, information which is unavailable from the previous studies in the literature."
" At z2:2 we are exclusively interested inthe orm of high-mass end of the local black-hole mass function. and consequentis. in this section we use the Adj,—ec and MayLose relations to derive two independent estimates of the local black-hole mass function for earlv-tvpe galaxies."," At $z\simeq 2$ we are exclusively interested inthe form of high-mass end of the local black-hole mass function, and consequently, in this section we use the $M_{bh}-\sigma$ and $M_{bh}-L_{bulge}$ relations to derive two independent estimates of the local black-hole mass function for early-type galaxies."
 Seth et al. (, Seth et al. (
2003) recently performed. a detailed analysis of the stellar-velocity. dispersion function of 9000. carly-type ealaxies drawn from the SDSS.,2003) recently performed a detailed analysis of the stellar-velocity dispersion function of 9000 early-type galaxies drawn from the SDSS.
 As part of this analysis Seth et al., As part of this analysis Seth et al.
" provide a fitting formula for the dispersion function. similar to a Schechter function in form. which is governed by four free parameters (6,.0,.0.:)."," provide a fitting formula for the dispersion function, similar to a Schechter function in form, which is governed by four free parameters $(\phi_{\star}, \sigma_{\star},\alpha,\beta)$."
 As the starting point [or our caleulation we adopt the best. fit determined: by Seth et al., As the starting point for our calculation we adopt the best fit determined by Seth et al.
 to the observed. dispersion function. including nmieasurement errors. which is described by the parameter values (0.002. Mpc SS km +. 6.5. L.S)," to the observed dispersion function, including measurement errors, which is described by the parameter values (0.002 $^{-3}$ , 88 km $^{-1}$ , 6.5, 1.8)."
 Secondly. we," Secondly, we"
Equation 2..,Equation \ref{eqn:multigauss}.
" Phe CAIB Uuetuations can also be expressed in spherical harmonics as where the e;,, are complex and can be written If the temperature field is a multivariate: Gaussian then the real ancl imaginary parts of the (e;,, should be mutually independent and Gaussian distributed (Bardeenet 1986).", The CMB fluctuations can also be expressed in spherical harmonics as where the $a_{\ell m}$ are complex and can be written If the temperature field is a multivariate Gaussian then the real and imaginary parts of the $a_{\ell m}$ should be mutually independent and Gaussian distributed \cite{bbks86}.
". Furthermore. the phases yr, should be random."," Furthermore, the phases $\varphi_{\ell m}$ should be random."
 1n our analyses we look for non.Gaussianity in the WALAP data in both real and harmonie spaces., In our analyses we look for non–Gaussianity in the WMAP data in both real and harmonic spaces.
 Both spaces allows us to probe a range of length scales., Both spaces allows us to probe a range of length scales.
 Harmonic space has the advantage of being more condensed., Harmonic space has the advantage of being more condensed.
 Whereas. the advantage of studsing in real space is that we can casily navigate contaminated regions or focus on a specific area in the CM skv.," Whereas, the advantage of studying in real space is that we can easily navigate contaminated regions or focus on a specific area in the CMB sky."
 The exact nature of the data used is fully described in Section 4.., The exact nature of the data used is fully described in Section \ref{sec:data}.
 In this Section. we outline the procedures. used. later to examine whether the WALAP data is strictly. multivariate Gaussian.," In this Section, we outline the procedures used later to examine whether the WMAP data is strictly multivariate Gaussian."
 Generally. when looking for signs of non-normalitv. it helps to have an idea of the form it should take.," Generally, when looking for signs of non-normality, it helps to have an idea of the form it should take."
 The wide variety of possible sources of nonCGaussianity means that there is no unique form of alternative distribution to seek., The wide variety of possible sources of non–Gaussianity means that there is no unique form of alternative distribution to seek.
 Fortunately. this is quite a common problem in multivariate analyses. where real data does not adhere to any specific alternative mocel.," Fortunately, this is quite a common problem in multivariate analyses, where real data does not adhere to any specific alternative model."
 This is unsurprising as there are very alternative models for which he entire hierarchy. of multivariate distributions is fully specified., This is unsurprising as there are very alternative models for which the entire hierarchy of multivariate distributions is fully specified.
 For this reason statisticians tend to apply not just one test statistic to the data. but a battery of complementaty oocedures.," For this reason statisticians tend to apply not just one test statistic to the data, but a battery of complementaty procedures."
 The assorted procedures. will have dillering sensitivities to the shape of the distribution., The assorted procedures will have differing sensitivities to the shape of the distribution.
 There is also a need to augment these tests with analyses of subsets of he data., There is also a need to augment these tests with analyses of subsets of the data.
 Testing the form Equation 2. in all its generality is Clearly impossible as it requires an infinite number of tests., Testing the form Equation \ref{eqn:multigauss} in all its generality is clearly impossible as it requires an infinite number of tests.
 In. practice. various simplified. approaches tend o be implemented: in particular. the one-point marginal distributions of the variates are often. studied.," In practice, various simplified approaches tend to be implemented: in particular, the one-point marginal distributions of the variates are often studied."
 Marginal normality does not imply joint normality. although the ack of multivariate. normality is often. reflected. in. the marginal distributions.," Marginal normality does not imply joint normality, although the lack of multivariate normality is often reflected in the marginal distributions."
 A further advantage of examining he mareinal distributions is that they are computationally ess intense and more intuitive (i.e. easier to interpret) and hus more instructive., A further advantage of examining the marginal distributions is that they are computationally less intense and more intuitive (i.e. easier to interpret) and thus more instructive.
 The procedures we apply to the data are described in hiree subsections., The procedures we apply to the data are described in three subsections.
 In subsection 3.1... we outline univariate echniques for assessing marginal normality.," In subsection \ref{sec:univariate}, , we outline univariate techniques for assessing marginal normality."
 In. subsection J. multivariate techniques for evaluating joint. normality are sketched out.," In subsection \ref{sec:multivariate}, multivariate techniques for evaluating joint normality are sketched out."
 Lastly. we illustrate a procedure. that evaluates the degree to which the regression of cach variate on all others is linear in subsection 3.3..," Lastly, we illustrate a procedure that evaluates the degree to which the regression of each variate on all others is linear in subsection \ref{sec:linear}. ."
 Vhroughout these subsections. we shall denote the members of the 7th variate. xi. bv ο where j=L....n.," Throughout these subsections, we shall denote the members of the $i$ th variate, $\mathv{x}_i$, by $x_{ij}$ where $j=1,\dots,n$."
 In subsection 3.2... we shall use the notation ΠΩ”...rpj). and similarly xi=GrasopespY.," In subsection \ref{sec:multivariate}, we shall use the notation $\mathv{x}_j=(x_{1j}, x_{2j}, \dots, x_{pj})'$ and similarly $\mathv{x}_k=(x_{1k}, x_{2k}, \dots, x_{pk})'$ ."
 We wish to emphasize the distinction between x; and Xx;=Gripe...rinY.," We wish to emphasize the distinction between $\mathv{x}_j$ and $\mathv{x}_i=(x_{i1},
x_{i2}, \dots, x_{in})'$."
 The evaluation of marginal normality of the data is based on wellknown tests of univariate normality., The evaluation of marginal normality of the data is based on well–known tests of univariate normality.
 Phe marginal distributions we study. correspond to the distribution for each individual variate., The marginal distributions we study correspond to the distribution for each individual variate.
 Fhis is simply the distribution of members of the specified variate. ignoring all other members from the dataset.," This is simply the distribution of members of the specified variate, ignoring all other members from the data–set."
" For example. in the bivariate case. the marginal probability distribution of the variate X, is given by The parallel expression for the marginal distributions corresponding to larger values of the dimensionalitv p can easily be developed."," For example, in the bivariate case, the marginal probability distribution of the variate $\mathv{x}_1$ is given by The parallel expression for the marginal distributions corresponding to larger values of the dimensionality $p$ can easily be developed."
 We outline four techniques that. probe univariate normality. of P(x;)y the skewness and. kurtosis cocllicients: 1)Agostino’s omnibus test: and a shifted:power transformation test.," We outline four techniques that probe univariate normality of ${\cal
P}(\mathv{x}_i)$: the skewness and kurtosis coefficients; D'Agostino's omnibus test; and a shifted–power transformation test."
 In this subsection. we shall suppress the { indices when referring to the datamembers such that oj will refer to the individual members of x;.," In this subsection, we shall suppress the $i$ indices when referring to the data–members such that $x_j$ will refer to the individual members of $\mathv{x}_i$ ."
" ‘The classic tests of normality is by means of evaluating the sample skewness V/b, and kurtosis b» coellicients.", The classic tests of normality is by means of evaluating the sample skewness $\sqrt{b_1}$ and kurtosis $b_2$ coefficients.
 LE we let the sample mean and the sample variance be wr and s? respectively. and define to be the 7/ moment about. the mean.," If we let the sample mean and the sample variance be $\bar{x}$ and $S^2$ respectively, and define to be the $r^{\mathrm th}$ moment about the mean."
 EThen. the skewness is given by and the kurtosis by The expected value of the skewness (Mardia1980). is Lere. and throughout the rest of the paper. we use the notation (60) for the expected value ofthe statistic Q and a(Q) to the signify the square root of the variance. aboutthis value.," Then the skewness is given by and the kurtosis by The expected value of the skewness \cite{m80} is and the square root of the variance of this quantity is Here, and throughout the rest of the paper, we use the notation $E(Q)$ for the expected value ofthe statistic $Q$ and $\sigma(Q)$ to the signify the square root of the variance aboutthis value."
 These values assume € is applied to aGaussian dataset., These values assume $Q$ is applied to aGaussian data–set.
 The same quantities for the kurtosis are, The same quantities for the kurtosis \cite{m80} are
The strong morphology segregation observed iu rich clusters of galaxies (Dressler. d980) testifies the fundamental role plawed bx the enviromment on the evolution of galaxies.,"The strong morphology segregation observed in rich clusters of galaxies (Dressler, 1980) testifies the fundamental role played by the environment on the evolution of galaxies."
 Which plysical imechanisuis are responsible for such transformations is however still matter of debate., Which physical mechanisms are responsible for such transformations is however still matter of debate.
 Several processes uuelt alter the evolution of cluster galaxies., Several processes might alter the evolution of cluster galaxies.
 Some of them refer to the interaction of the galaxies with the intracluster uedimu (Comm Cott. 1972) and others account for the effects of eravitational interactions produced by the gravitational potential of the cluster (Alerritt. 1983) or bv galaxy-ealaxy interactions (Moore et al.," Some of them refer to the interaction of the galaxies with the intracluster medium (Gunn Gott, 1972) and others account for the effects of gravitational interactions produced by the gravitational potential of the cluster (Merritt, 1983) or by galaxy-galaxy interactions (Moore et al."
 1996. 1908. 1999).," 1996, 1998, 1999)."
 All these mechanisnis can produce stroug perturbations iu the ealaxy morphology with the formation of tidal tails. dynamical disturbances which appear as asviunietries in the rotation curves (Dale ct al.," All these mechanisms can produce strong perturbations in the galaxy morphology with the formation of tidal tails, dynamical disturbances which appear as asymmetries in the rotation curves (Dale et al."
 2001) and significant eas removal (Cüovanelli Tavnes 1985: Vallu Jog 1990)., 2001) and significant gas removal (Giovanelli Haynes 1985; Valluri Jog 1990).
 Some of these processes are expected to produce changes in the star formation rates of ealaxics in clusters., Some of these processes are expected to produce changes in the star formation rates of galaxies in clusters.
 Several studies have addressed the issue of the influence of the cluster euvironmoent on the SFR of disk ealaxies. however no agreement has been established so far: whereas some authors proposed simular or even enhanced star formation im cluster spirals than iu the field (Donas ct al.," Several studies have addressed the issue of the influence of the cluster environment on the SFR of disk galaxies, however no agreement has been established so far: whereas some authors proposed similar or even enhanced star formation in cluster spirals than in the field (Donas et al."
 1990. 1995: Moss Whittle 1993. Cavazzi Coutursi 1991: Moss et al.," 1990, 1995; Moss Whittle 1993, Gavazzi Contursi 1994; Moss et al."
 1998: Ciavazzi et al., 1998; Gavazzi et al.
 1998: Moss Whittle 2000). some others chui quenched SFRs in cluster spirals (I&eunicutt 1983: Balogh ot al.," 1998; Moss Whittle 2000), some others claim quenched SFRs in cluster spirals (Kennicutt 1983; Balogh et al."
 1998: IHashinoto et al., 1998; Hashimoto et al.
 1998)., 1998).
 This discrepancy could arise from non-uniformity of the adopted methods (UV vs. Ta vs. Ομ data) or from real differences in the studied. clusters (Virgo. Coma. Abell 1367. clusters from Las Campanas Redshift Survey. clusters at +20.18). Iu particular. an cuhanced fraction of spirals with circiununuclear Πα enüsson was found in the highest density regious of some nearby clusters (Moss et al.," This discrepancy could arise from non-uniformity of the adopted methods (UV vs. $\alpha$ vs. ] data) or from real differences in the studied clusters (Virgo, Coma, Abell 1367, clusters from Las Campanas Redshift Survey, clusters at $z > 0.18$ In particular, an enhanced fraction of spirals with circumnuclear $\alpha$ emission was found in the highest density regions of some nearby clusters (Moss et al."
 1998:, 1998;
"slightly, such that the total spectral width covered by both filters is comparable.","slightly, such that the total spectral width covered by both filters is comparable."
" Details on the data reduction, object detection, and selection of NB-excess sources can be found in ?.."," Details on the data reduction, object detection, and selection of NB-excess sources can be found in \citet{Kurk2004a}."
" Candidate Ha emitters were selected following the same procedure as described above, in particular, the X=2 curve for the 11138-262 data follows closely the dotted curve in refha;elect(seeFig.66in?))."," Candidate $\alpha$ emitters were selected following the same procedure as described above, in particular, the $\Sigma=2$ curve for the 1138-262 data follows closely the dotted curve in \\ref{ha_select} (see 6 in \citealt{Kurk2004a}) )."
ToallowarobustcomparisonweonlyselectedN. excesssourceswithrest— frameEWo>30AA., To allow a robust comparison we only selected NB-excess sources with rest-frame $EW_{\rm 0}>30$.
". There are 38 sources with X>2 and EW>30AA,, of which 3 are within the Lya halo of the radio galaxy, and thus not counted as part of the overall large-scale structure."," There are 38 sources with $\Sigma>2$ and $>30$, of which 3 are within the $\alpha$ halo of the radio galaxy, and thus not counted as part of the overall large-scale structure."
 Spectroscopy of 9 of the brightest candidate Ha emitters are presented in ?.., Spectroscopy of 9 of the brightest candidate $\alpha$ emitters are presented in \citet{Kurk2004b}.
 All are confirmed to have an emission line and 3 of which are confirmed to be Ha due to the presence of [Nit]., All are confirmed to have an emission line and 3 of which are confirmed to be $\alpha$ due to the presence of ].
" The other 6 objects with just one emission line are also clustered in velocity space near the redshift of the radio galaxy, which strongly suggests that they are located in a large-scale structure associated with the radio galaxy."," The other 6 objects with just one emission line are also clustered in velocity space near the redshift of the radio galaxy, which strongly suggests that they are located in a large-scale structure associated with the radio galaxy."
" Field Ha emitting candidates at z—2.2 were selected using HHAWE-I narrow-band images of the GOODS-S, UDS and COSMOS fields obtained through programs 081.A-0932(A) (see ? for details), and 083.A-0826(A) (PII. P. Best)."," Field $\alpha$ emitting candidates at $z\sim2.2$ were selected using HAWK-I narrow-band images of the GOODS-S, UDS and COSMOS fields obtained through programs 081.A-0932(A) (see \citealt{Hayes2010} for details), and 083.A-0826(A) (P.I. P. Best)."
" The image of the GOODS-S was obtained with the HHAWK-I narrow-band filter which is wwide, so captures eemitted from galaxies at 2.178«z2.207."," The image of the GOODS-S was obtained with the HAWK-I narrow-band filter which is wide, so captures emitted from galaxies at $2.178<z<2.207$."
" The images of COSMOS and the UDS were obtained with the nnarrow-band filter which is wwide, so they capture eemitted from galaxies at 2.213<z«2.259."," The images of COSMOS and the UDS were obtained with the narrow-band filter which is wide, so they capture emitted from galaxies at $2.213<z<2.259$."
 All 3 fields have a similar field-of-view of aarcmin?., All 3 fields have a similar field-of-view of $^2$.
 The data were reduced using MVM in the same manner as described above., The data were reduced using MVM in the same manner as described above.
 Images taken in poor seeing were not included in the final combined dataset., Images taken in poor seeing were not included in the final combined dataset.
 The UDS was not observed in the Ks by HAWK-I so the DR8 K image from UKIDSS was used instead., The UDS was not observed in the $Ks$ by HAWK-I so the DR8 $K$ image from UKIDSS was used instead.
" Final exposure times, seeing full width at half maximum (FWHM), and resulting 1σ limiting magnitudes are given in reftab:fieldobs."," Final exposure times, seeing full width at half maximum (FWHM), and resulting $\sigma$ limiting magnitudes are given in \\ref{tab:field_obs}."
.T heK/Ks-banddatawere fluxcalibratedusingthe publiclyavailablecat ," The $K/Ks-$ band data were flux calibrated using the publicly available catalogues of \citet{Retzlaff2010}, Simpson et (2010 in prep.)"
, and \citet{Ilbert2009}.
NBcoloursof starsinthe fields.," The NB data were flux calibrated using standard stars, and further adjustments made after examining the $Ks-NB$ colours of stars in the fields."
"T heresolutiono ftheimageswerehomogenisedtomat cht resolutionimages, andcoloursweremeasuredinaperturesthatwere3timestheseeingF W ProrGOODS—S,2.4"" fortheUDS, and2.1"" forCOSMOS)"," The resolution of the images were homogenised to match that of the lowest-resolution images, and colours were measured in apertures that were 3 times the seeing FWHM for GOODS-S, for the UDS, and for COSMOS)."
 Inall3control f ieldsapproximatel y90960fthelis , In all 3 control fields approximately of the light from a point source lies within these apertures.
"Objects were detected using in double image mode, using a weighted NB image as the detecting image."," Objects were detected using in double image mode, using a weighted $NB$ image as the detecting image."
 The sameSEXTRACTOR parameter file used for the ddata was also used for the control field data to ensure the source detection was as similar as possible., The same parameter file used for the data was also used for the control field data to ensure the source detection was as similar as possible.
 ggalaxies were selected as objects with EWo>30À aand X>2 (see reffield-selection))., galaxies were selected as objects with $EW_{\rm 0}>30$ and $\Sigma>2$ (see \\ref{field-selection}) ).
 We ensured that the selection of candidates was identical to the ffield by selecting galaxies to the same X=2 limit., We ensured that the selection of candidates was identical to the field by selecting galaxies to the same $\Sigma=2$ limit.
 The COSMOS data are the shallowest data therefore only sources up to the X=2 depth of the COSMOS data were selected., The COSMOS data are the shallowest data therefore only sources up to the $\Sigma=2$ depth of the COSMOS data were selected.
" This selection resulted in 18 ggalaxies in GOODS-S, 13 in the UDS, and 31 in COSMOS."," This selection resulted in 18 galaxies in GOODS-S, 13 in the UDS, and 31 in COSMOS."
 All candidates were checked by eye to ensure they were real sources., All candidates were checked by eye to ensure they were real sources.
 The total combined area of all 3 control fields is aarcmin?., The total combined area of all 3 control fields is $^{2}$ .
uodel.,model.
 The results are eiven in Table 3.., The results are given in Table \ref{tab:rrspec}.
 In all fits we wave frozen the metallicity of the thermal coupoucuts to hat of the general cluster cussion., In all fits we have frozen the metallicity of the thermal components to that of the general cluster emission.
 If allowed to vary. the uctallicity is uncoustramed.," If allowed to vary, the metallicity is unconstrained."
 The single tempcrature gas uodel. while a decent fit to the data. docs not reproduce he excess cluission iu the data near OS keV (see Figure 9)).," The single temperature gas model, while a decent fit to the data, does not reproduce the excess emission in the data near 0.8 keV (see Figure \ref{fig:rrspec}) )."
 The two-temperature eas model is the best fit o the data., The two-temperature gas model is the best fit to the data.
 Oue of the compoucuts is consistent with the overall cluster temperature. while the other componcut las a temperature of 20.9 keV aud accounts for tle soft excess seen iu the data.," One of the components is consistent with the overall cluster temperature, while the other component has a temperature of $\approx$ 0.9 keV and accounts for the soft excess seen in the data."
 The gas | powerlaw imodel is a worse fit. with the powerlaw component replacing the totter gas componcut.," The gas $+$ powerlaw model is a worse fit, with the powerlaw component replacing the hotter gas component."
 The two-teniperature gas model is the most appropriate. as it would include both a lower cluperature eas associated with the relic as well as hotter cluster enission seen il projection.," The two-temperature gas model is the most appropriate, as it would include both a lower temperature gas associated with the relic as well as hotter cluster emission seen in projection."
 We have presented the data of the galaxy cluster. Abell 13., We have presented the data of the galaxy cluster Abell 13.
 The X-ray ciission is centered ou galaxy TD of Sleeetal.(2001). the brightest cluster galaxy.," The X-ray emission is centered on galaxy H of \citet{srm+01}, the brightest cluster galaxy."
 The spatial structure near the core is complicated. with a bright tail of excess enission to the west of galaxy II. aud a fainter excess to the north.," The spatial structure near the core is complicated, with a bright tail of excess emission to the west of galaxy H, and a fainter excess to the north."
 No point source is associated with galaxy II and no cooling flow is found., No point source is associated with galaxy H and no cooling flow is found.
 We also found excess N-ray endssou associated with the secoud-brightest elliptical galaxy. EF. in the cluster.," We also found excess X-ray emission associated with the second-brightest elliptical galaxy, F, in the cluster."
 The mass of gas associated with galaxy F is consistent with other bright elliptical galaxies iu cluster environments (Vikhlininetal.2001:Sun2007).," The mass of gas associated with galaxy F is consistent with other bright elliptical galaxies in cluster environments \citep{vmf+01,sjf+07}."
. Iu general. there is an extension of the X-ray cussion to the northwest with respect to the X-ray peak. coiucideut with ealaxy IT. Finally. no excess or significant deficit iu enission is found coimcident with the uuusual radio relic in Abell 13. although the spectral properties of the eas in this region are different from those in other areas of the cluster.," In general, there is an extension of the X-ray emission to the northwest with respect to the X-ray peak coincident with galaxy H. Finally, no excess or significant deficit in emission is found coincident with the unusual radio relic in Abell 13, although the spectral properties of the gas in this region are different from those in other areas of the cluster."
 We found a soft excess in the spectrum of eas coincident with the radio relic. whichcan be well fit by a 0.9 keV thermal model. suggesting the presence of cooler gas there.," We found a soft excess in the spectrum of gas coincident with the radio relic, whichcan be well fit by a $\approx$ 0.9 keV thermal model, suggesting the presence of cooler gas there."
 The redshift distribution of the galaxies in Abell 13 can be described by a bimodal distribution (Faddaetal.1996)., The redshift distribution of the galaxies in Abell 13 can be described by a bimodal distribution \citep{fgg+96}.
. The subchuups are not distinguishable iu their location on the sky., The subclumps are not distinguishable in their location on the sky.
 Tuterestingly. the two brightest ealaxies. IT aud. E. are associated with differeut peaks in the redshift distribution aud rather than being associated with the peak velocity. either of the cluster as a whole or of either subclip. the velocities of galaxies IT aud F are on the very edees of the velocity distribution.," Interestingly, the two brightest galaxies, H and F, are associated with different peaks in the redshift distribution and rather than being associated with the peak velocity, either of the cluster as a whole or of either subclump, the velocities of galaxies H and F are on the very edges of the velocity distribution."
 Galaxy F is on the lower οσο for the low-velocity subcluster aud. galaxy IT on the higher edge of the ligh-velocity subcluster., Galaxy F is on the lower edge for the low-velocity subcluster and galaxy H on the higher edge of the high-velocity subcluster.
 This. combined with the complicated structure of the ταν cussion. poiuts to au ongoiug morecr within the cluster and that galaxies ID aud F have a significant velocity relative to the general cluster distribution.," This, combined with the complicated structure of the X-ray emission, points to an ongoing merger within the cluster and that galaxies H and F have a significant velocity relative to the general cluster distribution."
 The spatial distribution of the ealaxics πι Abell 13 shows a nortlwest-southeast extension. paralleling the X-ray distribution. but exteudiug| to larger scales.," The spatial distribution of the galaxies in Abell 13 shows a northwest-southeast extension, paralleling the X-ray distribution, but extending to larger scales."
 Iu the following discussion. we assume that the N-rav eas associated with the radio relic has two compoucuts. a hot (6 keV) gas component due to the general cluster enission. and a cooler (0.9 keV) eas component local to he relie.," In the following discussion, we assume that the X-ray gas associated with the radio relic has two components, a hot (6 keV) gas component due to the general cluster emission, and a cooler (0.9 keV) gas component local to the relic."
 This model is based. on our best-fit spectrin oundins 77.., This model is based on our best-fit spectrum found in \ref{sec:rad}. .
 We asstune that the two componcuts are in pressure equilibrium., We assume that the two components are in pressure equilibrium.
" From Sleeetal. (2001).. the αι pressure of he radio relie in Abcll 13 is 1.1.«104? dyne cm7,"," From \citet{srm+01}, , the minimum pressure of the radio relic in Abell 13 is $1.4 \times 10^{-12}$ dyne $^{-2}$."
 The cstimate of the thermal pressure of the hot X-ray gas was very rough aud based solely ou scaling the xoperties of Abell 85 since no imaging observation of Abell 13 was available., Their estimate of the thermal pressure of the hot X-ray gas was very rough and based solely on scaling the properties of Abell 85 since no imaging observation of Abell 13 was available.
 OurChendra data allow for a nore accurate estimate of the thermal pressure at the ocation of the relic., Our data allow for a more accurate estimate of the thermal pressure at the location of the relic.
 However. due to an unfortunate dlacement of the cluster on the CCD. the only regious which extended out to a distance larger than the relic were in the direction of the relic aud to the south.," However, due to an unfortunate placement of the cluster on the CCD, the only regions which extended out to a distance larger than the relic were in the direction of the relic and to the south."
 Since he southern edge of the cluster shows a steep eradieut in the ταν cuuission. we did not consider this region in determining the surface brightucss distribution.," Since the southern edge of the cluster shows a steep gradient in the X-ray emission, we did not consider this region in determining the surface brightness distribution."
 Iustead. we determined the surface brightness of the cluster in a conical region extending frou the cluster center out o LAL and eucompassing position angles of 312° ucasured from the north to the cast.," Instead, we determined the surface brightness of the cluster in a conical region extending from the cluster center out to 1 and encompassing position angles of $\arcdeg$ measured from the north to the east."
 This region includes he radio relic aud associated cool gas which will affect our results: however. while not a perfect measure of the xoperties of the N-rav eas. if represents a siguificaut inroveinienut on the estimates of Sleeetal.i(2001).," This region includes the radio relic and associated cool gas which will affect our results; however, while not a perfect measure of the properties of the X-ray gas, it represents a significant improvement on the estimates of \citet{srm+01}."
. A arecr field of view observation would be better suited o address the surface brightuess aud electron density distributions of the full cluster emission., A larger field of view observation would be better suited to address the surface brightness and electron density distributions of the full cluster emission.
 Using the surface briehtuess distribution. we estimate he electron. deusity distribution of the hot componcut ollowiug the spherical deprojection method of Eris (1983)...," Using the surface brightness distribution, we estimate the electron density distribution of the hot component following the spherical deprojection method of \citet{kcc83}. ."
 We assiuneda single temperature of the rot gasof 6.0 keV and used the thermal model, We assumeda single temperature of the hot gasof 6.0 keV and used the thermal model
Solar-like oscillations in non-relativistic stars are thought o be stochastically excited by turbulent convection in he outer lavers of the star.,Solar-like oscillations in non-relativistic stars are thought to be stochastically excited by turbulent convection in the outer layers of the star.
 Vheorctical investigations (Goldreich&οσον 1977: Balmilorth 1992: Gioldreich.Alurray&Ixumar 1994: Samacdi&Coupil 2001)) identified wo major mechanisms by which oscillations are. driven: he first source term is due to the turbulent Reynolds stress and is thus related to the Revnolels stress. tensor: he second mechanism stems from the advection of Eulerian entropy Bluctuations by turbulent motions (the entropy source term)., Theoretical investigations \citealt{Goldreich1977}; ; \citealt{Balmforth1992}; ; \citealt*{Goldreich1994}; \citealt{Samadi2001a}) ) identified two major mechanisms by which oscillations are driven: the first source term is due to the turbulent Reynolds stress and is thus related to the Reynolds stress tensor; the second mechanism stems from the advection of Eulerian entropy fluctuations by turbulent motions (the entropy source term).
 A generalized formalism by Samaci&Goupil(2001) for radial modes. which attempts at. combining roth mechanisms in a consistent. forced wave equation. was [απο extended by Belkacemetal.(2008) o the case of non-radial modes.," A generalized formalism by \citet{Samadi2001a} for radial modes, which attempts at combining both mechanisms in a consistent forced wave equation, was further extended by \citet{Belkacem2008} to the case of non-radial modes."
 Vheoretically computed excitation rates for radial solar p modes are in good. agreement with observational data from the GOLF instrument on board the SOLO spacecraft. (Samadictal.2003a.b::: Belkacemetal. 2006a.h2).," Theoretically computed excitation rates for radial solar p modes are in good agreement with observational data from the GOLF instrument on board the SOHO spacecraft \citealt{Samadi2003a,Samadi2003b}; \citealt{Belkacem2006a,Belkacem2006b}) )."
. Furthermore. radial solar p-mode peak heights were shown to be reasonably well modelled (Chaplinetal. 2005: Loudek 2006)).," Furthermore, radial solar p-mode peak heights were shown to be reasonably well modelled \citealt{Chaplin2005}; \citealt{Houdek2006}) )."
 However. in case of solar ο modes the entropy. source term is negligible (Belkacemοἱal.2000). and quantitative estimates of mode amplitudes differ from cach other by orders of magnitude. depending mainly on the assumect ecldy-time correlation function. that is on how the turbulent eclelies are time-correlatect (Ixumar.Quataert&Baheall 1906: Belkacemetal. 2009: see in particular the discussion in Xppourchauxetal. 2010)).," However, in case of solar g modes the entropy source term is negligible \citep{Belkacem2009} and quantitative estimates of mode amplitudes differ from each other by orders of magnitude, depending mainly on the assumed eddy-time correlation function, that is on how the turbulent eddies are time-correlated \citealt*{Kumar1996}; \citealt{Belkacem2009}; see in particular the discussion in \citealt{Appourchaux2010}) )."
 Since 1e detection. of solar ο modes still remains controversial. ieoretical predictions of their amplitudes are of utmost importance.," Since the detection of solar g modes still remains controversial, theoretical predictions of their amplitudes are of utmost importance."
 In the present paper. we consider an independent riving mechanism which we expect to be relevant for solar η. solar-like & modes. especially in case of stars larger iin the Sun that are nearby strong. gravitational wave sources: excitation by external gravitational waves.," In the present paper, we consider an independent driving mechanism which we expect to be relevant for solar and solar-like g modes, especially in case of stars larger than the Sun that are nearby strong gravitational wave sources: excitation by external gravitational waves."
 Due to amping rates. /—1 and /=2 solar ο modes are the most probable candidates fordetection (Belkacemetal.2009).," Due to damping rates, $l=1$ and $l=2$ solar g modes are the most probable candidates for detection \citep{Belkacem2009}."
 As is shown below. this mechanism is able to excite quacrupolar cigenmocdes (/= 2) and thus it may influence the future detection of solar and solar-like & modes.," As is shown below, this mechanism is able to excite quadrupolar eigenmodes $l=2$ ) and thus it may influence the future detection of solar and solar-like g modes."
 The excitation of normal modes by gravitational waves has already been studied for relativistic stars., The excitation of normal modes by gravitational waves has already been studied for relativistic stars.
 Dhese stars were shown to have a family of normal modes that are directly associatedwiththe curvature of spacetime. called eravitational-wave modes (wo modes). and. that have no Newtonian analogue in non-relativistic stars," These stars were shown to have a family of normal modes that are directly associatedwiththe curvature of space–time, called gravitational-wave modes (w modes), and that have no Newtonian analogue in non-relativistic stars"
the (νι 5254) maser line. (1.1). (2.2) and (3.3) inversion transitions. radio recombination lines LIG2ea and 11690 and LCCCN (32).,"the $_{1,6}$ $_{2,3}$ ) maser line, (1,1), (2,2) and (3,3) inversion transitions, radio recombination lines $\alpha$ and $\alpha$ and HCCCN (3–2)."
 masers are an important. signpost of unusual astrophysical— conditions such as outllows ancl shocked. gas., masers are an important signpost of unusual astrophysical conditions such as outflows and shocked gas.
 They are known to occur in both high and low-mass star forming regions (eg. ForsterandCaswell1999:Claussenetal. 1996))," They are known to occur in both high and low-mass star forming regions (eg. \citealt{forster99,claussen96}) ),"
. late Metvpe stars (Dickinson1976).. planetary nebulae (Mirandaetal. 2001)... Mira. variables Barnes 1979).. Asvimptotie Ciant Branch stars al.1996) and the centres of active galaxies (Claussenctal. 1984).," late M-type stars \citep{dickinson76}, planetary nebulae \citep{miranda01}, , Mira variables \citep{hinkle79}, Asymptotic Giant Branch stars \citep{barlow96} and the centres of active galaxies \citep{claussen84}."
. Phe majority of currently known masers are Found towards regions of star formation within our Galaxy., The majority of currently known masers are found towards regions of star formation within our Galaxy.
 Llowever. an untargeted survey is required to determine the relative occurrence of bright masers with other tvpes of astrophysical objects.," However, an untargeted survey is required to determine the relative occurrence of bright masers with other types of astrophysical objects."
 Phere is some evidence that within star forming regions. masers may be observable at very carly stages (eg. ForsterandCaswell 2000))," There is some evidence that within star forming regions, masers may be observable at very early stages (eg. \citealt{forster00}) )"
 of evolution., of evolution.
 There is evidence that methanol masers may also be visible at these carly stages (eg. Walshetal. 1998))., There is evidence that methanol masers may also be visible at these early stages (eg. \citealt{walsh98}) ).
 The relative occurrence of these two masers can be assessed through the two untarected surveys: the Methanol. Alultibeam Survey (Greenetal.2009) and LOPS. described. here.," The relative occurrence of these two masers can be assessed through the two untargeted surveys: the Methanol Multibeam Survey \citep{green09} and HOPS, described here."
 Recent rescarch has focused on determining parallax distances to masers (eg. Imaiοἱal. 2007))., Recent research has focused on determining parallax distances to masers (eg. \citealt{imai07}) ).
 A lone term goal is to accurately establish the dimensions of our Galaxy using bright maser sources throughout the Galaxy (ee. Reicletal. 20093)., A long term goal is to accurately establish the dimensions of our Galaxy using bright maser sources throughout the Galaxy (eg. \citealt{reid09}) ).
 It is hoped that the bright masers discovered in LOPS may be used for such distance —determinations., It is hoped that the bright masers discovered in HOPS may be used for such distance determinations.
 Thermal emission in our Galaxy typically traces righ density (~LOtem: Evans 1999)) gas.," Thermal emission in our Galaxy typically traces high density $\sim
10^4 {\rm cm}^{-3}$; \citealt{evans99}) ) gas."
 Properties of spectral line emission can be used to understand the yhvsical conditions of the gas., Properties of spectral line emission can be used to understand the physical conditions of the gas.
 Lower order (J.Ix) inversion ransitions of commonly display hyvperfine structure (ος. Lo&Townes 1983))," Lower order (J,K) inversion transitions of commonly display hyperfine structure (eg. \citealt{ho83}) )"
 which can be used to calculate the optical depth of the emission. to reduce confusing factors of optically thick emission. when trving to measure the otal column density of eas.," which can be used to calculate the optical depth of the emission, to reduce confusing factors of optically thick emission, when trying to measure the total column density of gas."
 Multiple inversion transitions ol are also commonly seen. which can be used to estimate the temperature of the excited gas.," Multiple inversion transitions of are also commonly seen, which can be used to estimate the temperature of the excited gas."
 These inversion ransitions occur within a few CGllz of each other. making it possible to observe them: simultaneously. with the same elescope and similar setup.," These inversion transitions occur within a few GHz of each other, making it possible to observe them simultaneously with the same telescope and similar setup."
 This greatly. eliminates sources of uncertainty when comparing multiple transitions. which in turn makes interpretation more reliable.," This greatly eliminates sources of uncertainty when comparing multiple transitions, which in turn makes interpretation more reliable."
 In cold. dense regions most molecules ancl ions are known to freeze-out. onto dust. grains. with CO. CS and being classic examples (Bergin&Langer1997:al.2002:Bereinet 2002).," In cold, dense regions most molecules and ions are known to freeze-out onto dust grains, with CO, CS and $^+$ being classic examples \citep{bergin97,willacy98,caselli99,tafalla02,bergin02}."
 Phus. such species are unreliable tracers of the densest and. coldest. regions in our Galaxy.," Thus, such species are unreliable tracers of the densest and coldest regions in our Galaxy."
 CO also has a low elfective critical density ~QU) ecm (Evans1990) making it a better tracer of dilluse gas ancl susceptible to becoming optically thick in dense regions., CO also has a low effective critical density $\sim10^2$ $^{-3}$ \citep{evans99} making it a better tracer of diffuse gas and susceptible to becoming optically thick in dense regions.
 CO. CS and — are also known to be good tracers of outllow activity in star forming regions.," CO, CS and $^+$ are also known to be good tracers of outflow activity in star forming regions."
 Therefore a Galactic map of emission in these tracers can be dillieult to interpret in terms of identifving quiescent gas., Therefore a Galactic map of emission in these tracers can be difficult to interpret in terms of identifying quiescent gas.
 On the other hand. appears to be more robust than CO. €S or against [reeze-out. onto cust grains (eg. Aikawaetal.— 20013).," On the other hand, appears to be more robust than CO, CS or $^+$ against freeze-out onto dust grains (eg. \citealt{aikawa01}) )."
 This means that can be used to probe the colder denser regions., This means that can be used to probe the colder denser regions.
 In addition to this. Nol which is found. under conditions very similar toNlI4.. is known to avoid regions of outllow (ce. Walshetal.2007a)).," In addition to this, $_2$ $^+$, which is found under conditions very similar to, is known to avoid regions of outflow (eg. \citealt{walsh07a}) ),"
 tracing only the quiescent gas on Large scales., tracing only the quiescent gas on large scales.
 “Vhis makes à very reliable tracer of cool. dense gas.," This makes a very reliable tracer of cool, dense gas."
 Combined with the information derived from hyperfine structure and multiple inversion transitions. can be used to reliably characterise the dense. quiescent gas component of the ISM.," Combined with the information derived from hyperfine structure and multiple inversion transitions, can be used to reliably characterise the dense, quiescent gas component of the ISM."
 In this paper. we concentrate on describing the survey as well as presenting global properties of the maser data.," In this paper, we concentrate on describing the survey as well as presenting global properties of the maser data."
 In Paper LE (Purcell et al., In Paper II (Purcell et al.
 2011. preparation) we will esent the NIL; (1.1) and (2.2) data and cloud. catalogue as well as describing the emission finding algorithm.," 2011, ) we will present the $_3$ (1,1) and (2,2) data and cloud catalogue as well as describing the emission finding algorithm."
 Paper LL (Longmore at al., Paper III (Longmore at al.
 2011. preparation) will detail the hermal line fitting routines used on the NIL; cata and will present the physical properties of the unresolved clouds in the catalogue.," 2011, ) will detail the thermal line fitting routines used on the $_3$ data and will present the physical properties of the unresolved clouds in the catalogue."
 Results from. other spectral lines will »* reported in later papers., Results from other spectral lines will be reported in later papers.
 In follow-up. work. we will accurately measure the positions of the masers using he Australia Telescope C'ompact Array (ATCA).," In follow-up work, we will accurately measure the positions of the masers using the Australia Telescope Compact Array (ATCA)."
 The Australia Telescope National Facility Mopra telescope is a 22mm antenna located. kkm outside the town of Coonabarrabran in New South Wales. Australia.," The Australia Telescope National Facility Mopra telescope is a m antenna located km outside the town of Coonabarrabran in New South Wales, Australia."
 Ht is at an elevation of 850 metres above sea level and at a Latitude of 31 south., It is at an elevation of 850 metres above sea level and at a latitude of $^\circ$ south.
 The receivers use Indium Phosphide Ligh Electron Mobility “Transistor (LIEAIT) Monolithic Microwave Integrated Circuits (MMICS) as the amplifving elements., The receivers use Indium Phosphide High Electron Mobility Transistor (HEMT) Monolithic Microwave Integrated Circuits (MMICs) as the amplifying elements.
 The receiver systems require no tuning ancl include noise diodes for svstem temperature determination., The receiver systems require no tuning and include noise diodes for system temperature determination.
 The nominal operating range for the 12mmim receiver is [rom 16 to GOGlLIz. however we have found the receiver still performs well at frequencies as high as GCGlIIz.," The nominal operating range for the mm receiver is from 16 to GHz, however we have found the receiver still performs well at frequencies as high as GHz."
 Further general information on Mopra can be found. in Urquhartet.al. (2010)., Further general information on Mopra can be found in \citet{urquhart10}.
. The Mopra Spectrometer (MODPS) is the digital filterband backend used for the observations., The Mopra Spectrometer (MOPS) is the digital filterband backend used for the observations.
 It comprises an GCllz total bandwidth. split into four overlapping intermediate frequencies (IPs). cach with a width of ολ].," It comprises an GHz total bandwidth, split into four overlapping intermediate frequencies (IFs), each with a width of GHz."
 MOPS can be conligured in two modes. either. broadband: mode. where the full S.3€:CLIz is available. or in zoom. mocle. which was used for LOPS.," MOPS can be configured in two modes, either broadband mode, where the full GHz is available, or in zoom mode, which was used for HOPS."
 In. zoom mocdoe. cach of the [our IPs contains four zoom bands of MMIIz cach. available to take spectra.," In zoom mode, each of the four IFs contains four zoom bands of MHz each, available to take spectra."
 The positioning of the four MMIIz zoom bands within each Ue is highly exible. which allows the user to observe virtually any of up to our spectral windows within cach αν IE.," The positioning of the four MHz zoom bands within each IF is highly flexible, which allows the user to observe virtually any of up to four spectral windows within each GHz IF."
 Thus. it is possible to simultaneously observe up to 16 spectral windows hroughout the full S.3€CGlIIz: bandwidth.," Thus, it is possible to simultaneously observe up to 16 spectral windows throughout the full GHz bandwidth."
 Each spectral window consists of 4096 channels. which is equivalent to a xucbiwidth aad velocity resolution of and at AHCGIIz. orancl at GGlI. respectively.," Each spectral window consists of 4096 channels, which is equivalent to a bandwidth and velocity resolution of and at GHz, orand at GHz, respectively."
 Fable 1. gives the details, Table \ref{tab1} gives the details
for flat cosmologies.,for flat cosmologies.
 The solid points indicate the clusters from the sample studied here., The solid points indicate the clusters from the sample studied here.
 The clusters from the 160SD survey for which X-ray temperatures have been determined (see $33.1) are indicated by blue points., The clusters from the 160SD survey for which X-ray temperatures have been determined (see 3.1) are indicated by blue points.
 To extend the range in X-ray luminosity we also show measurements for the massive clusters that were studied in ?.., To extend the range in X-ray luminosity we also show measurements for the massive clusters that were studied in \cite{Hoekstra07}.
 We note. that ?— used bolometric X-ray luminosities from ?.. whereas here we use the restframe luminosities in the 0.1—2.4 keV band (which are a factor ~4 smaller).," We note, that \cite{Hoekstra07} used bolometric X-ray luminosities from \cite{Horner01}, whereas here we use the restframe luminosities in the $0.1-2.4$ keV band (which are a factor $\sim 4$ smaller)."
 We use the mass estimates from ? which used new photometric redshift distributions. which were based on much larger data sets (?)..," We use the mass estimates from \cite{Mahdavi08}
 which used new photometric redshift distributions, which were based on much larger data sets \citep{Ilbert06}."
 The agreement in lensing masses is good in the regions where the two samples overlap., The agreement in lensing masses is good in the regions where the two samples overlap.
 However. the scatter in the mass-luminosity relation is substantial (both for the clusters studied here as well as the more massive clusters).," However, the scatter in the mass-luminosity relation is substantial (both for the clusters studied here as well as the more massive clusters)."
 We examined whether some of the scatter could be due to the uncertainty in the position of the cluster center. but find no difference when comparing results for clusters with different levels of confidence in the identification of the BCG (see Qpcc in Table ..," We examined whether some of the scatter could be due to the uncertainty in the position of the cluster center, but find no difference when comparing results for clusters with different levels of confidence in the identification of the BCG (see $Q_{\rm
 BCG}$ in Table \ref{tabsample}."
" It is worth noting. however. that we do not find any massive clusters (>107M..) that are X-ray faint (i.e.. Ly<10"" erg/s). implying that the dispersion in the M—Ly relation is relatively well-behaved."," It is worth noting, however, that we do not find any massive clusters $(>10^{14}M_{\odot})$ that are X-ray faint (i.e., $L_X<10^{44}$ erg/s), implying that the dispersion in the $M-L_X$ relation is relatively well-behaved."
 ? studied a sample of 63 X-ray bright clusters and derived masses under theassumption of hydrostatie equilibrium., \cite{Reiprich02} studied a sample of 63 X-ray bright clusters and derived masses under theassumption of hydrostatic equilibrium.
 This sample of clusters spans a similar range in Ly as our combined sample., This sample of clusters spans a similar range in $L_X$ as our combined sample.
 We converted their measurements for άκου to Masao using the mass-concentration relation given by Eqn., We converted their measurements for $M_{500}$ to $M_{2500}$ using the mass-concentration relation given by Eqn.
 10 and show the results in Figure 6 (small grey points)., \ref{mcrel} and show the results in Figure \ref{mlx} (small grey points).
 The agreement with our findings is very good., The agreement with our findings is very good.
 A more quantitative comparison is presented in $33.3., A more quantitative comparison is presented in 3.3.
 ? have shown that there is à good relation between the cluster richness and the mass (and various proxies) for massive clusters., \cite{Yee03} have shown that there is a good relation between the cluster richness and the mass (and various proxies) for massive clusters.
 Our study extends to lower masses and as is shown in Figures 6bb and ο Nasco correlates well with both the X-ray luminosity and lensingmass!., Our study extends to lower masses and as is shown in Figures \ref{mlx}b b and c $N_{2500}$ correlates well with both the X-ray luminosity and lensing.
. Similar results have been obtained using SDSS cluster samples (e.g..??)..," Similar results have been obtained using SDSS cluster samples \citep[e.g.,][]{Rykoff08a,Johnston07}."
 We assumed that N»soo scales as the mass M»soo. which is a reasonable assumption if the galaxies trace the density profile.," We assumed that $N_{2500}$ scales as the mass $M_{2500}$, which is a reasonable assumption if the galaxies trace the density profile."
 This choice. however. does not affect our conclusions.," This choice, however, does not affect our conclusions."
 The results agree well in the regions where the two samples overlap., The results agree well in the regions where the two samples overlap.
 The correlation between Λήξης and Mosoo is tighter than that of Λο and Ly., The correlation between $N_{2500}$ and $M_{2500}$ is tighter than that of $N_{2500}$ and $L_X$.
 The former is less sensitive to the projections along the line-of-sight (either substructures or an overall elongation of the cluster). because both the galaxy counts and the lensing mass are derived from projected measurements.," The former is less sensitive to the projections along the line-of-sight (either substructures or an overall elongation of the cluster), because both the galaxy counts and the lensing mass are derived from projected measurements."
 The X-ray results provide a different probe of the distribution of baryons. which is expected to lead to additional seatter.," The X-ray results provide a different probe of the distribution of baryons, which is expected to lead to additional scatter."
 Furthermore. some of the scatter may be caused by unknown contributions by AGNs.," Furthermore, some of the scatter may be caused by unknown contributions by AGNs."
 The importance of AGN can be evaluated using a combination of deeper. high spatial resolution (<537) X-ray and radio imaging.," The importance of AGN can be evaluated using a combination of deeper, high spatial resolution $(\lesssim 5'')$ X-ray and radio imaging."
 Such X-ray observations would also provide estimates for the temperature of the X-ray gas (which isa better measure of the cluster mass), Such X-ray observations would also provide estimates for the temperature of the X-ray gas (which isa better measure of the cluster mass).
 Unfortunately. such data exist for only five of the low-mass clusters.," Unfortunately, such data exist for only five of the low-mass clusters."
 For these. X-ray temperatures have been derived. which are listed Table 3..," For these, X-ray temperatures have been derived, which are listed Table \ref{tabtx}. ."
 For the massive clusters we use the values from ?. that were used in ?.., For the massive clusters we use the values from \cite{Horner01} that were used in \cite{Hoekstra07}. .
" We note that all clusters follow a tight L,—7, relation.", We note that all clusters follow a tight $L_x-T_x$ relation.
Furthermore. the cross derivatives. F;;=wewith /4j ae defined: acas In Eqs. €,"Furthermore, the cross derivatives, $F_{ij} = \frac{\partial F_{j}}{\partial{x_{i}}}$with $i \ne j$, are defined as In Eqs. ("
"CA.D)-CA.2). the terms 7). £5. Ps. Ty. Ts. Tu. and T; have following forms: fj,=(25-απx). T.=ίαασ]. r=(8) = (r=(eS) 22). and 7;=(:- 21) ","A.1)-(A.2), the terms$T_{1}$ , $T_{2}$ $T_{3}$ , $T_{4}$,$T_{5}$ , $T_{6}$ , and $T_{7}$ have following forms: $T_{1} = \left(2x_{3}^2 - x_{1}^2 - x_{2}^2\right)$, $T_{2} = \left(x_{1}^2 - x_{2}^2\right)$, $T_{3} = \left(\frac{v_{D}} {2r^2} - \frac{v_{r}} {r^3}\right)$, $T_{4} = \left(\frac{w_{D}} {r^2} - \frac{2w_{r}} {r^3}\right)$, $T_{5} = \left(\frac{v_{DD}} {2r^2} - \frac{2v_{D}} {r^3} + \frac{3v_{r}} {r^4}\right)$ $T_{6} = \left(\frac{w_{DD}} {r^2} - \frac{4w_{D}} {r^3} + \frac{6w_{r}} {r^4}\right)$ , and $T_{7} = \left(\frac {1} {r} - \frac {x_{1}^2} {r^3}\right)$ ."
Here typ.Vip. and wp aregiven by which represent the derivatives of Πρ. ," Here$u_{DD}$ ,$v_{DD}$ and $w_{DD}$ aregiven by which represent the derivatives of $u_{D}$ "
If the velocity change in Fig.,If the velocity change in Fig.
 6 is due to Increasing inass. it also implies that the radia ejection velocities will decrease with time.," 6 is due to increasing mass, it also implies that the radial ejection velocities will decrease with time."
 Since constant radial velocities were assumed iu the calculation of the ages of cach triplet. this wouk meni that the initial ejection velocities were hieher than preseutle measured aud ages wouk be slightly lower than calculated.," Since constant radial velocities were assumed in the calculation of the ages of each triplet, this would mean that the initial ejection velocities were higher than presently measured and ages would be slightly lower than calculated."
 Before this cau be confirmed. redshifts for objects LO aud 12 wil need to be measured.," Before this can be confirmed, redshifts for objects 10 and 12 will need to be measured."
 NatlikaraudDas(1980)/ sugeeeested several tess that wuelt be used to check their theory., \citet{nar80b} suggested several tests that might be used to check their theory.
 Oue of these predicted. that the redsuft depeudence of outiele inasses in a QSO should show up iu its linunosity., One of these predicted that the redshift dependence of particle masses in a QSO should show up in its luminosity.
 Tf the svuchrotrei mechanisni is yonsible for enmuüsson. the Lhnuuinositv should relate to the QSO intrinsic redshift (zo) by the uctor (1 | zo) (NarnlikarandDas—1980)..," If the synchrotron mechanism is responsible for emission, the luminosity should relate to the QSO intrinsic redshift $_{\rm Q}$ ) by the factor (1 + $_{\rm Q})^{2}$ \citep{nar80b}."
 Iun Fig., In Fig.
 5(a) the heavy solid iue shows how 1ο apparent imagnitude would be expected. to change with redshift for this relaion., 5(a) the heavy solid line shows how the apparent magnitude would be expected to change with redshift for this relation.
 The values rave been normalized Το Ζωα 1.2., The values have been normalized to $_{\rm mean}$ = 1.2.
 Although 1¢ luminosity appears to increase more slowly iu predicted by Narlikay aud Das this cau be explained., Although the luminosity appears to increase more slowly than predicted by Narlikar and Das this can be explained.
 If it is assmmed that their uuniunositv change occurs in the ceveloping jiebulositv or “host” galaxy. here would prestmably need to be a correction required to compensae for the simultaneous dinuiue of the ceutra qlauar.," If it is assumed that their luminosity change occurs in the developing nebulosity or ""host"" galaxy, there would presumably need to be a correction required to compensate for the simultaneous dimming of the central quasar."
 In t1ο Narlikar and Das theory the maunuer in which uass changes with age (intrinsic redshift) i:4. also eiveni, In the Narlikar and Das theory the manner in which mass changes with age (intrinsic redshift) is also given.
 As an exiuuple they point out that ifa QSO of redshift zi = 2 is seen in association with a galaxy of redshift ze: = 0.005. he ithe electron mass iu the latter is predicted to be 3 times the[um electron mass in the former.," As an example they point out that if a QSO of redshift $_{\rm Q}$ = 2 is seen in association with a galaxy of redshift $_{\rm G}$ = 0.005, then the electron mass in the latter is predicted to be 3 times the electron mass in the former."
 Iu he case of NGC LOGS. if he redshift iu Fig.," In the case of NGC 1068, if the redshift in Fig."
 6 is Doppler reated. as concluded in Paper L aud is decrease with tine Is ¢ue to mereasing lass. a crude estimate of the inagnitude of this mass Πιοοπο uieht be determined by makiug the siuple assuuiptio- that mouentuni is conserved.," 6 is Doppler related, as concluded in Paper I, and its decrease with time is due to increasing mass, a crude estimate of the magnitude of this mass increase might be determined by making the simple assumption that momentum is conserved."
 After makine this assuniptiou. the relative mass as a function of time is plotted in Fig.," After making this assumption, the relative mass as a function of time is plotted in Fig."
 7 aud. as above. is 1n reasonable agreement with what has been predicted by the Narlikur aud Das model.," 7 and, as above, is in reasonable agreement with what has been predicted by the Narlikar and Das model."
"aims of HerCULES are the development of the diagnostic use of the gas cooling lines in local (U)LIRGs, and establishing a local benchmark for observations of high-z galaxies with the Atacama Large Millimeter Array.","aims of HerCULES are the development of the diagnostic use of the gas cooling lines in local (U)LIRGs, and establishing a local benchmark for observations of $z$ galaxies with the Atacama Large Millimeter Array."
" In addition, since the FTS yields full spectra, any other luminous emission lines detected (e.g., of H2O) will be available for study."," In addition, since the FTS yields full spectra, any other luminous emission lines detected (e.g., of $\HtO$ ) will be available for study."
" ? have shown that X-ray excitation of the gas (e.g., by an AGN) and UV irradiation by young massive stars produce very different luminosity distributions over the CO rotational lines."," \citet{SpaansMeijerink08} have shown that X-ray excitation of the gas (e.g., by an AGN) and UV irradiation by young massive stars produce very different luminosity distributions over the CO rotational lines."
" Physically, the difference arises because X-rays penetrate a larger column density of gas than UV photons, and are less effective in dissociating the molecules."," Physically, the difference arises because X-rays penetrate a larger column density of gas than UV photons, and are less effective in dissociating the molecules."
" In addition, while the gas heating efficiency in a photon dominated region (PDR) is less than 1%, in X-ray dominated regions (XDRs) this efficiency is 10-50%."," In addition, while the gas heating efficiency in a photon dominated region (PDR) is less than $1\%$, in X-ray dominated regions (XDRs) this efficiency is $10{-}50\%$."
" As a result, for comparable irradiated energies, XDRs tend to have larger column densities of warmer molecular gas than PDRs, and will produce much more luminous emission in the high-J CO lines."," As a result, for comparable irradiated energies, XDRs tend to have larger column densities of warmer molecular gas than PDRs, and will produce much more luminous emission in the $J$ CO lines."
" In contrast, PDRs"," In contrast, PDRs"
A-ray Haring behaviour but for which an ee counterpart is νο to be determined (e.g.Sgueraetal.ma2006).,X-ray flaring behaviour but for which an optical counterpart is yet to be determined \citep[e.g.][]{SFXTsguera}.
. The outburst: histories of some of m SENXTS have been shown to display periodic emission. interpreted as the orbital period of the compact object in the system.," The outburst histories of some of the known SFXTs have been shown to display periodic emission, interpreted as the orbital period of the compact object in the system."
 These periods have a large range from a few days. such as LG. 4514 with a period of clays (Jain.Paul.&Dutta 2009).. to fractions of a vear: LOR 5952 having an orbital period of 165ddays. (Romanoct 2009a).," These periods have a large range from a few days, such as IGR $-$ 4514 with a period of days \citep{Jain2009}, to fractions of a year; IGR $-$ 5952 having an orbital period of days \citep{Romano2009}."
. In some cases an X-ray. pulse period has also been discovered (e.gIG1tJ11215.5952.LGR.016465.4507.ICal. 2001).. lcading to the firm identification of the compact object as a neutron star.," In some cases an X-ray pulse period has also been discovered \citep[e.g IGR J11215$-$5952, IGR J16465$-$4507, IGR J18483$-$0311, AX J1841.0-0536,][]{2007ATel..999....1S,2006A&A...453..133W,2007A&A...467..249S,Bamba2001}, leading to the firm identification of the compact object as a neutron star."
 Several mechanisms to explain the observed outbursts in SEXT systems have been proposed., Several mechanisms to explain the observed outbursts in SFXT systems have been proposed.
 In general these use either a structure in the supergiant stellar wind or magnetic elfects on the accretion Dow., In general these use either a structure in the supergiant stellar wind or magnetic effects on the accretion flow.
 See Sidoli09b) and the references therein for a fuller review of the current. proposed. outburst mechanisms., See \citet{Sidolimechrev} and the references therein for a fuller review of the current proposed outburst mechanisms.
 Further analysis of outbursts from NTE 302 found. [ares lasting between 30 minutes anc 3 hours and variations in absorption column na seen between bursts (Smithetal.2006:Sguera2. die," Further analysis of outbursts from XTE $-$ 302 found flares lasting between 30 minutes and 3 hours and variations in absorption column density seen between bursts \citep{SmithJ173912006,SgueraJ173912005}."
dPhe X-ray spectrum of NTE 21739. 302 has been mi in all levels of emission (Sicolietal.2008.. Siclolietal. 2O009a.. Romanoetal. 2009b.. Blavetal. 2008.. Bozzoetal. 2010)).," The X-ray spectrum of XTE $-$ 302 has been well studied in all levels of emission \citealt{Sidoli2008_OutOut}, \citealt{Sidoli2008_InOut}, \citealt{Romano2009b}, \citealt{Blay2008}, \citealt{Bozzo2010XMM}) )."
 XTE J11739. 302 is often well fit with powerlaw spectra and exhibits variations in photon-index and absorption column densities between outburst and. quiescence. further. traits Of SENT behaviour.," XTE $-$ 302 is often well fit with powerlaw spectra and exhibits variations in photon-index and absorption column densities between outburst and quiescence, further traits of SFXT behaviour."
 Broadband spectra have also shown a possible cut-olf at ~13kkeV ancl indications of absorption features at 230 kkeV and. ~GOkkeVY. Phese spectral shapes are characteristic of an accreting neutron star (Negueruelaetal., Broadband spectra have also shown a possible cut-off at $\sim$ keV and indications of absorption features at $\sim$ keV and $\sim$ keV. These spectral shapes are characteristic of an accreting neutron star \citep{SFXTneguer}.
 2006a).. AXE 302 is located NE 2 degrees from the Galactic centre (IX. 17:39:i Dee -30:20:37.6 (Smithοἱal. 3))) and as a resul a large amoun of coverage from the Galactic Centre Bulge Alonitoring Campaign (Ixuulkersetal.2007). and Galactic Contre Deep Exposure (Winkleretal.1999)..," XTE $-$ 302 is located approximately 2 degrees from the Galactic centre (RA 17:39:11.6, Dec -30:20:37.6 \citep{SmithJ173912006}) ) and as a result has a large amount of coverage from the Galactic Centre Bulge Monitoring Campaign \citep{KuulkersGBMP} and Galactic Centre Deep Exposure \citep{INTEGRALcore}."
 In this repor new analysis undertaken on the full IBIS data set is presented., In this report new analysis undertaken on the full /IBIS data set is presented.
 Section 2 reports the data selection used anc Section 3 the analysis performed. in the form of periodicity analysis in Section 3.1 and outburst identification in Section 3.2.," Section 2 reports the data selection used and Section 3 the analysis performed, in the form of periodicity analysis in Section 3.1 and outburst identification in Section 3.2."
 A discussion of the results and. their physica interpretation is given in Section 4. followed by conclusions in Section 5," A discussion of the results and their physical interpretation is given in Section 4, followed by conclusions in Section 5."
 Our analysis used data from archives covering 552671.7 through 554763.6. an initial data set o£ 714.4 MMs. Using the Ollline Science Analysis (OSA: Golelwurmetal. 2003)). ντ software. an 15 - OkkeV Science. Window (ScW) LBIS/ISGRI lishteurve was generated. (Birdetal.2010)..," Our analysis used data from archives covering 52671.7 through 54763.6, an initial data set of $\sim$ Ms. Using the Offline Science Analysis (OSA; \citealt{GoldwurmOSA}) ) v7.0 software an 18 - keV Science Window (ScW) IBIS/ISGRI lightcurve was generated \citep{TonyCat4}."
 Nominally a ScW represents an ~2000ss observation., Nominally a ScW represents an $\sim$ s observation.
 The lighteurve was filtered as to disregard SeWs with exposure times of less than ss and/or source locations separated from the telescope pointing axis by more than 12 degrees., The lightcurve was filtered as to disregard ScWs with exposure times of less than s and/or source locations separated from the telescope pointing axis by more than 12 degrees.
 As a result. SeWs with large errors on the flux. determination are removed. which is desirable when data is to be used for Lomb-Scargle analvsis. see Section 3.1.," As a result, ScWs with large errors on the flux determination are removed, which is desirable when data is to be used for Lomb-Scargle analysis, see Section 3.1."
 D result is an IS - 6OkkeY lighteurve. spanning 552698.2n to 554763.6 with a total cllective exposure of MSs. ATE 41739302 also has coverage in the All Sky Monitor (ASAD) ane Burst Alert Felescope (BAT) public databases.," The result is an 18 - keV lightcurve, spanning 52698.2 to 54763.6 with a total effective exposure of $\sim$ Ms. XTE $-$ 302 also has coverage in the All Sky Monitor (ASM) and Burst Alert Telescope (BAT) public databases."
 After filtering for poor quality data μασ». short exposures ancl low coded aperture fractions. the {BAT lighteurve covers 553413 through to 555208 with an elfective exposure time of ~13.2 MINIs. The RAXTE//ASM. liehteurve spans MJD550088.1. to 555209.9. with a total exposure time of ~3.45 MIAIs. The filtered LBIS/ISGRIE lighteurve was tested for M signals by means of a Lomb-Scarele analysis (Lomb197€M Scargle1982) di," After filtering for poor quality data flags, short exposures and low coded aperture fractions, the /BAT lightcurve covers 53413 through to 55208 with an effective exposure time of $\sim$ Ms. The /ASM lightcurve spans 50088.1 to 55209.9, with a total exposure time of $\sim$ Ms. The filtered IBIS/ISGRI lightcurve was tested for periodic signals by means of a Lomb-Scargle analysis \citealt{LOMB1976}, , \citealt{SCARGLE1982}) )."
ede- Phe resulting periodogram contained seve peaks of high power., The resulting periodogram contained several peaks of relatively high power.
 To test for significance a Monte-Carlo. based. randomisation test. was performed. as outlined in Hilletal.," To test for significance a Monte-Carlo based randomisation test was performed, as outlined in \citet{AdamMCT}."
 The resulting. periodogram ancl calculated confidence levels are shown in Fig. 1.., The resulting periodogram and calculated confidence levels are shown in Fig. \ref{fig1}.
 As Fig., As Fig.
 1: shows. there are several peaks that appear to be above the significance level.," \ref{fig1} shows, there are several peaks that appear to be above the significance level."
 Llowever. as a binary system. containing this number of physical periodicities is unlikely. efforts were mace to explain the presence of cach signal.," However, as a binary system containing this number of physical periodicities is unlikely, efforts were made to explain the presence of each signal."
 The first οσο considered. was that. of the window function resulting from the large scale. non-uniform sampling of data within the EDIS/ISCHIUI lighteurve.," The first effect considered was that of the window function resulting from the large scale, non-uniform sampling of data within the IBIS/ISGRI lightcurve."
 To evaluate these effects. simulated. lighteurves with the same shape as the phase-folded. lightcurves (c.g. Fig. 2))," To evaluate these effects, simulated lightcurves with the same shape as the phase-folded lightcurves (e.g. Fig. \ref{fig2}) )"
 and identical data gaps to the real data were eenerated. subjected to Lomb-Scarele analysis and their periodograms produced.," and identical data gaps to the real data were generated, subjected to Lomb-Scargle analysis and their periodograms produced."
 Only a period of ddays reproduced patterns seen in the actual. periodogram: this being also the period with the maximum power in the data., Only a period of days reproduced patterns seen in the actual periodogram; this being also the period with the maximum power in the data.
 As seen in the lower panel of Fig., As seen in the lower panel of Fig.
 1. the periodogram. of this window function explains the large features at ~40 and ddavs as well as features at. shorter periods., \ref{fig1} the periodogram of this window function explains the large features at $\sim$ 40 and days as well as features at shorter periods.
 These short. period features are consistent with many. of the periods. above confidence in the real data and. are. interpreted as higher frequency components introduced. by. the sinusoidal profile of the phase-folded lighteurve., These short period features are consistent with many of the periods above confidence in the real data and are interpreted as higher frequency components introduced by the non-sinusoidal profile of the phase-folded lightcurve.
 While the low period. peaks do not entirely match hese seen in the real periodogram. this analysis was only vcrlormed for the time-averagecl phase-folded: lighteurve.," While the low period peaks do not entirely match these seen in the real periodogram, this analysis was only performed for the time-averaged phase-folded lightcurve."
 The analysis was repeated for three subsets of the data. and the peaks at short periods were found to vary between subsets (e.g. the lowest inset in Fig. 1)).," The analysis was repeated for three subsets of the data, and the peaks at short periods were found to vary between subsets (e.g. the lowest inset in Fig. \ref{fig1}) ),"
 most likely due o changes in observing pattern and frequency during he mission., most likely due to changes in observing pattern and frequency during the mission.
 Furthermore there are other possible sources of periocicity within this period range., Furthermore there are other possible sources of periodicity within this period range.
 The orbital period is ~3 days and during the core. program (Winkleretal.1999) was performing scans of thegalactic plane ancl galactic centre exposures in à very regimented manner., The orbital period is $\sim$ 3 days and during the core program \citep{INTEGRALcore} was performing scans of thegalactic plane and galactic centre exposures in a very regimented manner.
 As a result the shorter period significant signals are taken as being either directly related. to the dday signal. from other sources of artificial periodicity," As a result the shorter period significant signals are taken as being either directly related to the day signal, from other sources of artificial periodicity"
 (Staneketal.2003:[orthKawabataοἱal.2003).. (e.e..Iwamotoetal.2003)..," \citep{Stan03,Hjo03,Kaw03}, \cite[e.g.,][]{Iwa03}. \citep{N94,WLW95}."
" (Nomotoοἱal.1994:Weaver1995).. M10AL.vr! e,~1000—4000kms|. (e.g..Chevalier&Fransson2003).. (e.g..Li&Chevalier2003)."," $\Mdot \sim 10^{-5} \ml$ $v_{\rm w} \sim 1000-4000\kms$ \citep[e.g.,][]{CF03}. \citep[e.g.,][]{LC03}."
. atypical Mazzalietal.(2002). Derger.Kulkarni.&Chevalier(2002) , \citet{Maz02} \citet{BKC02} 
appears to decrease during the simulation. before reaching à minimum value of (17:3520.005 bevond which it decreases no further.,"appears to decrease during the simulation, before reaching a minimum value of $\langle 1/\beta \rangle \simeq 0.005$ beyond which it decreases no further."
 Such a decrease did not occur during the other elobal runs (€11-C/3). or in the calculation of a massive planet in a turbulent disc in paper LL," Such a decrease did not occur during the other global runs (G1-G3), or in the calculation of a massive planet in a turbulent disc in paper II."
 At the present time it is unclear whether the turbulent energy is being allected by the protoplanet in such a manner that the turbulent ονnamo is operating less cllicicnthy due to the ongoing planetary perturbation. or whether the decrease is the result ofa large but temporary fluctuation induced by inserting the planet instantaneously in the turbulent disc.," At the present time it is unclear whether the turbulent energy is being affected by the protoplanet in such a manner that the turbulent dynamo is operating less efficiently due to the ongoing planetary perturbation, or whether the decrease is the result of a large but temporary fluctuation induced by inserting the planet instantaneously in the turbulent disc."
 Phe former scenario is conceivable when one considers that the azimuthal domain in run G5 is w/2 so that the fluid. experiences the (strong) »erturbation of the planet four times more frequently than it would do in a full 2x disc., The former scenario is conceivable when one considers that the azimuthal domain in run G5 is $\pi/2$ so that the fluid experiences the (strong) perturbation of the planet four times more frequently than it would do in a full $2 \pi$ disc.
 Such a strong ancl freeuent »erturbation may be able to alfect the underlying dynamo in such a wav as to reduce the magnetic energy., Such a strong and frequent perturbation may be able to affect the underlying dynamo in such a way as to reduce the magnetic energy.
 The shearing (ox runs in general do not support this picture. but run Db4 did also show a reduction in magnetic energv once he planet was inserted.," The shearing box runs in general do not support this picture, but run Bb4 did also show a reduction in magnetic energy once the planet was inserted."
 Lt is interesting to note that the results of Winters. Balbus. Llawley (2003a) also contain a suggestion that the saturation level of the MED turbulence may be alfected by the presence of a giant. planet.," It is interesting to note that the results of Winters, Balbus, Hawley (2003a) also contain a suggestion that the saturation level of the MHD turbulence may be affected by the presence of a giant planet."
 We now consider the impact of the embedded protoplanets on the velocity field of the disc. and specifically the point at which. the horseshoe orbits are clearly established.," We now consider the impact of the embedded protoplanets on the velocity field of the disc, and specifically the point at which the horseshoe orbits are clearly established."
 Figure 26 shows the Iuid trajectories for two runs., Figure \ref{fig25} shows the fluid trajectories for two runs.
" “Phe upper panel corresponds to a global calculation in which a protoplanet with Adp/Ad,=103 is enibedded in a aminar clise (i.e. the protoplanet is the same as that in run G3).", The upper panel corresponds to a global calculation in which a protoplanet with $M_P/M_*=10^{-4}$ is embedded in a laminar disc (i.e. the protoplanet is the same as that in run G3).
 Close inspection of this plot shows that the horseshoe rajectories are established. in this case., Close inspection of this plot shows that the horseshoe trajectories are established in this case.
 The lower panel shows the velocity field for run G3. where it is apparent that he horseshoe trajectories are disrupted. by the turbulent velocity field.," The lower panel shows the velocity field for run G3, where it is apparent that the horseshoe trajectories are disrupted by the turbulent velocity field."
 En other words the turbulence determines the Μη trajectories in this case. and the gravitational potential due to the planet is unable to establish the horseshoe orbits in the coorbital zone that are obtained in a (αμα disc model.," In other words the turbulence determines the fluid trajectories in this case, and the gravitational potential due to the planet is unable to establish the horseshoe orbits in the coorbital zone that are obtained in a laminar disc model."
 Figure 27. shows the Uuicl trajectories for run 5., Figure \ref{fig26} shows the fluid trajectories for run G5.
 The jiorseshoes orbits in this case are very clearly defined. in agreement with the shearing box run Dad.," The horseshoes orbits in this case are very clearly defined, in agreement with the shearing box run Ba4."
 In this model he circulating region around the planet in the Hill sphere is also clearly visible., In this model the circulating region around the planet in the Hill sphere is also clearly visible.
 Phe etfect of the turbulence. on he velocity. field. in such a strongly perturbed model is essentially indiscernible in the near vicinity of the planet., The effect of the turbulence on the velocity field in such a strongly perturbed model is essentially indiscernible in the near vicinity of the planet.
 In paper Ll the circulating region arouncl the »otoplanet in the Lill sphere was found to be disrupted. and his was tentatively ascribed to magnetic braking caused by magnetic linkage between the cireumplanetary disc and the xotostellar clise (see also figure 20)).," In paper II the circulating region around the protoplanet in the Hill sphere was found to be disrupted, and this was tentatively ascribed to magnetic braking caused by magnetic linkage between the circumplanetary disc and the protostellar disc (see also figure \ref{fig19}) )."
 A similar situation was also found in the shearing box runs Da3 and Bat. where the usual circulating region in the Lill sphere was found to be absent.," A similar situation was also found in the shearing box runs Ba3 and Ba4, where the usual circulating region in the Hill sphere was found to be absent."
 In these runs the gravitational softening adopted was bc0.344. which is quite Large. and results in gas that accretes into the Lill sphere forming an ‘atmosphere’ that is largely pressure supported but with some angular momentum.," In these runs the gravitational softening adopted was $b \simeq 0.3H$, which is quite large, and results in gas that accretes into the Hill sphere forming an `atmosphere' that is largely pressure supported but with some angular momentum."
 In this case the removal of angular momentum will leac to a reduction of the spin of the atmosphere. as observed.," In this case the removal of angular momentum will lead to a reduction of the spin of the atmosphere, as observed."
 The softening used in run G5 was b=0.14. and so the formation ofa rotationally supported cireumplanetary disc is more pronounced.," The softening used in run G5 was $b = 0.1H$, and so the formation of a rotationally supported circumplanetary disc is more pronounced."
 Phe removal of angular momentum in this case should. allow further gas accretion into the Lll sphere without a modification of the rotation profile., The removal of angular momentum in this case should allow further gas accretion into the Hill sphere without a modification of the rotation profile.
 An AHID simulation should. therefore result in more material accreting into the planetary Lill sphere than occurs in a non magnetised run. if magnetic braking is indeed important.," An MHD simulation should therefore result in more material accreting into the planetary Hill sphere than occurs in a non magnetised run, if magnetic braking is indeed important."
 La order to test this we performed a simulation that was similar to run G5. except that magnetic fields were neglected. and ana viscosity was emploved with a—7.10. in a laminar disc model.," In order to test this we performed a simulation that was similar to run G5, except that magnetic fields were neglected and an $\alpha$ viscosity was employed with $\alpha=7 \times 10^{-3}$ in a laminar disc model."
 This laminar disc calculation was evolved. for identical amount. of time ⋅⊀to run G5 (/⋅ =1543.1401 -.anequivalent to 2 62.12planetary orbits). and resulted in less mass being aceretecd into the planetary Hill sphere than," This laminar disc calculation was evolved for an identical amount of time to run G5 $t=1543.14$ $\Omega^{-1}$, equivalent to $\simeq 62.12$ planetary orbits), and resulted in less mass being accreted into the planetary Hill sphere than"
1ionitoriug program was moved to the O.3-n Celestron telescope of the automated Abbey Ridge Observatory (Lane2007:Majaessetal.2008).. aud preseutlv collects a few hours of observation of Cepheids aud other suspected variables ou every available clear night —. amounting to several hundred hours of monitoring cach vear.,"monitoring program was moved to the 0.3-m Celestron telescope of the automated Abbey Ridge Observatory \citep{11,5}, and presently collects a few hours of observation of Cepheids and other suspected variables on every available clear night — amounting to several hundred hours of monitoring each year."
 Extended runs have also been mace for special objects.c.y.. primary unin for eclipsing svstenis aud siuall auplitude Cepheids.," Extended runs have also been made for special objects, primary minimum for eclipsing systems and small amplitude Cepheids."
 Some of the more interesting results frou such studies are presented here., Some of the more interesting results from such studies are presented here.
 This binay Cepheid of pulsation period P~9! is categorized as Type IT ou the basis of metallicity (ILuris1981).. but its rapid period changes also sugeest a first-crossing object (Turner2006).," This binary Cepheid of pulsation period $P \simeq 9^{\rm d}$ is categorized as Type II on the basis of metallicity \citep{7}, but its rapid period changes also suggest a first-crossing object \citep{2}."
. Reeular monitoring frou the ARO to establish its rate of period change revealed curious temporal changes in niean brightuess of the star that. while perhaps typical of Type ID variables (Templeton&TTenden 2009).. niv be linked to the stars unusual biuaritv: 35 AL. companion dn a close 1107.19 orbit (Παπάς&Welch 1989).," Regular monitoring from the ARO to establish its rate of period change revealed curious temporal changes in mean brightness of the star that, while perhaps typical of Type II variables \citep{8}, may be linked to the star's unusual binarity: 3–5 $M_{\odot}$ companion in a close $110^{\rm d}.49$ orbit \citep{6}. ."
. The light curve for six fines the present 95.157 pulsation period (Fig. 1))," The light curve for six times the present $9^{\rm d}.157$ pulsation period (Fig. \ref{ixcas}) ),"
.ie. ~£ P4 displays siunsoidal variatious sinilar to ucarside heating iu close binaries. and averaging of the same data (to remove the pulsational variations) phased to the 1107.19 orbital period suggests he possibility of eclipses iu the svstem.," $\sim \frac{1}{2} P_{\rm orb}$ , displays sinusoidal variations similar to nearside heating in close binaries, and averaging of the same data (to remove the pulsational variations) phased to the $110^{\rm d}.49$ orbital period suggests the possibility of eclipses in the system."
 Further work is uceded., Further work is needed.
 This category includes loug-period Cepheids like SV. Vul aud S Vul that exhibit raudoini fluctuations in pulsation period superposed upou the evolutionary trends in their OC diagrams (Fig. 2)).," This category includes long-period Cepheids like SV Vul and S Vul that exhibit random fluctuations in pulsation period superposed upon the evolutionary trends in their O–C diagrams (Fig. \ref{ocdiagrams}) ),"
 iu the case of SV Vul sueecsting a recent Increase in its rate of blueward evolution (Turucr&Derduikov 2001).. and recognized binaries Like X Cye (not illustrated).," in the case of SV Vul suggesting a recent increase in its rate of blueward evolution \citep{10}, and recognized binaries like X Cyg (not illustrated)."
 The ο€ data for SV Vul and ο Vul include analyses of AAVSO aud ARO observations for the stars., The O–C data for SV Vul and S Vul include analyses of AAVSO and ARO observations for the stars.
 Regular photoelectric mouitoring of Polaris is done at the DGO during nights of good. seciug in periods of the vear when photometric weatler occurs regularlv: typically unid-winter. spring. aud autuimu.," Regular photoelectric monitoring of Polaris is done at the BGO during nights of good seeing in periods of the year when photometric weather occurs regularly: typically mid-winter, spring, and autumn."
 The star is observed differentially relative to ΠΟ 5911 (spectral type À3 V. V= 5.56). and the maeuitucde scale is adjusted for παπα] aiy mass differeuces. between it and Polaris.," The star is observed differentially relative to HD 5914 (spectral type A3 V, $V = 5.86$ ), and the magnitude scale is adjusted for small air mass differences between it and Polaris."
 Typical sets of observations from 2003 to 2009 are shown iu Fig. 3..," Typical sets of observations from 2003 to 2009 are shown in Fig. \ref{polaris},"
 from which oue cau detect the eradual changes in phase resulting fromthe regular period increase of Polaris. as well as its eradually increasing pulsation amplitude.," from which one can detect the gradual changes in phase resulting fromthe regular period increase of Polaris, as well as its gradually increasing pulsation amplitude."
"Next, we relate the luminosity and the TIR luminosity.","Next, we relate the luminosity and the TIR luminosity."
" There is a good correlation between Laxarr and [τη (total infrared luminosity emitted by dust) in the FIS Bright Source Catalogue, the first primary catalogue from the Sky Survey (Yamamuraetal.2009),, as shown by Takeuchietal. (2010):: With this empirical formula, Drrr=(6.9+1.4)x10° for 881."," There is a good correlation between $L_\mathit{AKARI}$ and $L_\mathrm{TIR}$ (total infrared luminosity emitted by dust) in the FIS Bright Source Catalogue, the first primary catalogue from the All-Sky Survey \citep{yamamura10}, as shown by \citet{takeuchi10}: With this empirical formula, $L_\mathrm{TIR}=(6.9\pm 1.4)\times 10^{9}~\mathrm{L}_{\sun}$ for 81."
" This luminosity is divided by 4zD? to obtain the totalLo infrared flux, Fprr=1.7x107 erg cm? s-!,"," This luminosity is divided by $4\pi D^2$ to obtain the total infrared flux, $F_\mathrm{TIR}=1.7\times 10^{-8}$ erg $^{-2}$ $^{-1}$."
" Considering that the flux may be uncertain by 20 per cent, it would be appropriate to put a 20 per cent error for each flux or luminosity."," Considering that the flux may be uncertain by 20 per cent, it would be appropriate to put a 20 per cent error for each flux or luminosity."
 Suzukiet(2010) obtained smaller luminosity (4.0+0.2)x10? Le for the total FIR luminosity., \citet{suzuki10} obtained smaller luminosity $(4.0\pm 0.2)\times 10^9$ $_{\sun}$ for the total FIR luminosity.
" They consider contributions from two components: a cold component with a temperature of 22 K and a warm component with a temperature of 64 K, but these two components do not include the contribution from the emission at <20um."," They consider contributions from two components: a cold component with a temperature of 22 K and a warm component with a temperature of 64 K, but these two components do not include the contribution from the emission at $<20~\micron$."
 We can also check the consistency of the dust mass with the data., We can also check the consistency of the dust mass with the data.
" The total dust mass, Ma, is estimated as where &, is the mass absorption coefficient at a frequency v. dus"," The total dust mass, $M_\mathrm{d}$, is estimated as where $\kappa_\nu$ is the mass absorption coefficient at a frequency $\nu$."
tt mass (3.4x107 Mo) obtained by the observation (Bendoal. 2010)., t mass $3.4\times 10^7~\mathrm{M}_{\sun}$ ) obtained by the observation \citep{bendo10}.
". In reffig:radial,, we show the radial profiles in the three bands."," In \\ref{fig:radial}, we show the radial profiles in the three bands."
 The radial distances (r) are deprojected by considering the position angle (157 degree) and the inclination angle (57 degree) taken from NED., The radial distances $r$ ) are deprojected by considering the position angle (157 degree) and the inclination angle (57 degree) taken from NED.
 The inclination angles given in NED range from 55 to 63 degree., The inclination angles given in NED range from 55 to 63 degree.
 Therefore the deprojected radius is uncertain by 10 per cent., Therefore the deprojected radius is uncertain by 10 per cent.
" The radius from the centre is divided into bins with a width of 50 arcsec, and the intensities of the pixels contained in each radius bin is averaged to obtain the intensity as a function of radius."," The radius from the centre is divided into bins with a width of 50 arcsec, and the intensities of the pixels contained in each radius bin is averaged to obtain the intensity as a function of radius."
 The standard deviation for each radius bin is also shownbar., The standard deviation for each radius bin is also shown.
 The flux errors are much smaller than the standard deviations., The flux errors are much smaller than the standard deviations.
 The radial profile can be divided into two components: bulge (r«150 arcsec) and disc (r2150 arcsec)., The radial profile can be divided into two components: bulge $r<150$ arcsec) and disc $r\geq 150$ arcsec).
 The 65 and umbandimagesshowabrightbulgewheretheintensityincreasesrdptuéforet Ero elimgtboshorte," The 65 and $\micron$ band images show a bright bulge where the intensity increases rapidly as the radius decreases, while the 140 $\micron$ image does not have such a sharp rise."
r thatre 120 wma ayeküdesdieav esuchasharpr PI wavelengt," The difference in the `sharpness' of the bulge component among different bands is due to the high dust temperatures, which reflect an intense radiation from the active galactic nucleus (AGN) or from the stars with high surface density."
hHerschelSREdata(Sauvageetal.2010)., The bulge component is not prominent also in the long-wavelength SPIRE data \citep{sauvage10}.
". For the disk part, all the intensities in the three bands show a linear slope along the radius."," For the disk part, all the intensities in the three bands show a linear slope along the radius."
" The linear slope means that the dust emission is radially more extended than the stellar emission, whose radial profile drops exponentially (e.g.Baggett,&Ander- 1998)."," The linear slope means that the dust emission is radially more extended than the stellar emission, whose radial profile drops exponentially \citep[e.g.][]{baggett98}."
". Although Sauvageetal.(2010) apply a fitting with exponential or Gaussian function to the FIR radial profile of the data, the fit is not necessarily good."," Although \citet{sauvage10}
 apply a fitting with exponential or Gaussian function to the FIR radial profile of the data, the fit is not necessarily good."
 The dispersion shown by the bars is relatively large around r~400 arcsec., The dispersion shown by the bars is relatively large around $r\sim 400$ arcsec.
 This radius corresponds to the radial extent of the spiral arms., This radius corresponds to the radial extent of the spiral arms.
" To compare the radial extents in the three bands, we show the radial profiles of the colours, /,(140um)/1,(90um) and I,(65um)/1I,(90um) in reffig:radial.lr.."," To compare the radial extents in the three bands, we show the radial profiles of the colours, $I_\nu (140~\micron )/I_\nu (90~\micron)$ and $I_\nu (65~\micron )/I_\nu (90~\micron)$ in \\ref{fig:radial_clr}."
" The formerrisesastheradiusincreases, u"," The former rises as the radius increases, while the latter dose not show such a clear trend."
"mcolourandtheplateauo fthe140 um- umcolouraroundr-6arcsecarecausedbytherelativelyhightemperatureinth, "," The slight enhancement of the $65\,\micron -140\,\micron$ colour and the plateau of the $140\,\micron -90\,\micron$ colour around $r=6$ arcsec are caused by the relatively high temperature in the spiral arms."
"Because of the linear-slope behaviour of the radial profile, the radial extent is well defined by the intercepts on the c axis in Fig. 3.."," Because of the linear-slope behaviour of the radial profile, the radial extent is well defined by the intercepts on the $x$ axis in Fig. \ref{fig:radial}."
" The radial extents defined in this way are 706, 674, and 702 arcsec for the 65, 90, and 140 um images, respectively."," The radial extents defined in this way are 706, 674, and 702 arcsec for the 65, 90, and 140 $\micron$ images, respectively."
 The reason why the emission is more extended in umthan?in90 um can be the radial gradient of Tr., The reason why the emission is more extended in $\micron$ than in $\micron$ can be the radial gradient of $T_\mathrm{LG}$.
" The 140um—90ium colour is directly converted into Tra byassuming a functional form of v?B,(Tra) for the SED."," The $140\,\micron -90\,\micron$ colour is directly converted into $T_\mathrm{LG}$ byassuming a functional form of $\nu^2B_\nu (T_\mathrm{LG})$ for the SED."
" While umisclosertotheintensitymaximumoftheS ED,90 um at the Wien side is more sensitive to the temperature change, and decreases more rapidly outwards."," While $\micron$ is closer to the intensity maximum of the SED, $\micron$ at the Wien side is more sensitive to the temperature change, and decreases more rapidly outwards."
 A positive 160um—70uum colour gradient and a negative temperature gradient of cold dust component are also shown with the data by (2006)..," A positive $160\,\micron -70\,\micron$ colour gradient and a negative temperature gradient of cold dust component are also shown with the data by \citet{perez06}. ."
The Schechter function. conversely fits all (he range in Iuminosities and Table 4 reports the data that come out [rom the fitting procedure.,"The Schechter function, conversely fits all the range in luminosities and Table \ref{para_Schechter_SDSS_data} reports the data that come out from the fitting procedure."
 Table 4 also reports ων the value in magnitude where the Schechter function peaks: this value is defined when a—1. otherwise we leave the box blank.," Table \ref{para_Schechter_SDSS_data} also reports $M_{p,max} $, the value in magnitude where the Schechter function peaks; this value is defined when $\alpha \> -1 $, otherwise we leave the box blank."
[function used in Miller&Browning(2003a) both encourage solutions with nonpredefined components. when such components are warranted by the presence of unknown classes in the data.,"function used in \citet{pami} both encourage solutions with nonpredefined components, when such components are warranted by the presence of unknown classes in the data."
" In (1)). itis the Z4; term which provides the impetus for lorming these unknown classes. since (his term approaches its maxinimun value (2,=1) in the nonpredefined component Case,"," In \ref{newlik}) ), it is the $\beta_{c|j}$ term which provides the impetus for forming these unknown classes, since this term approaches its maximum value $\beta_{u|j}=1$ ) in the nonpredefined component case."
 The mixture modeling approach provides several inference capabilities when dealing with mixed labeled/unlabeled data sets ancl possibly unknown classes: 1) it allows one to inler whether or not a given sample belongs (ο one of the known classes: 2) it identifies purelv. unlabeled mixture components/clusters. which are reasonably treated as putative unknown classes or. al any rate. components of unknown classes: 3) conditioned on a known class hypothesis for a given sample. the model can infer from which known class (he sample originates (i.e.. (he usual classification inlerence capability).," The mixture modeling approach provides several inference capabilities when dealing with mixed labeled/unlabeled data sets and possibly unknown classes: 1) it allows one to infer whether or not a given sample belongs to one of the known classes; 2) it identifies purely unlabeled mixture components/clusters, which are reasonably treated as putative unknown classes or, at any rate, components of unknown classes; 3) conditioned on a known class hypothesis for a given sample, the model can infer from which known class the sample originates (i.e., the usual classification inference capability)."
" While the mixture modeling approach is naturally suited (to new class discovery. given mixed labeled/unlabeled data. neural network (NN) classifiers do not appear to be predisposed to making these inferences,"," While the mixture modeling approach is naturally suited to new class discovery given mixed labeled/unlabeled data, neural network (NN) classifiers do not appear to be predisposed to making these inferences."
 Neural networks are generally trained using a purely supervised approach. with class labels provided [ον every example in the training set.," Neural networks are generally trained using a purely supervised approach, with class labels provided for every example in the training set."
 Thus. in general. unlabeled samples plav no role in the training — given a nüxed labeled/unlabeled data set. the NN training will discaid all the examples [rom unknown classes. or. perhaps worse. erroneotsly impute ancl use known class labels for this unknown class data.," Thus, in general, unlabeled samples play no role in the training – given a mixed labeled/unlabeled data set, the NN training will discard all the examples from unknown classes, or, perhaps worse, erroneously impute and use known class labels for this unknown class data."
 Aecordinglv. the neural network is only explicitly trained to discriminate between the known classes — it is nol trained to distinguish known from unknown classes.," Accordingly, the neural network is only explicitly trained to discriminate between the known classes – it is not trained to distinguish known from unknown classes."
 While it (hus appears that neural networks do nol possess any class discovery inference capability. we next suggest an approach that give NNs at least a weak form of this capability.," While it thus appears that neural networks do not possess any class discovery inference capability, we next suggest an approach that give NNs at least a weak form of this capability."
 The neural network algorithm we used is a basic backpropagation algorithim available wilh the WEIXA machine learning package Witten&Frank.(2000)., The neural network algorithm we used is a basic backpropagation algorithm available with the WEKA machine learning package \citet{weka}.
. We used the default configuration consisting of a three laver network (input. hidden. aud output).," We used the default configuration consisting of a three layer network (input, hidden, and output)."
 The number of input nodes was N;. one node for each input feature.," The number of input nodes was $N_{i}$, one node for each input feature."
" There were N,. output nodes. one for each known class."," There were $N_{c}$ output nodes, one for each known class."
" The number of hidden nodes was caleulated according to Ny,=(NV;-:N.)/2.", The number of hidden nodes was calculated according to $N_{h} = (N_{i} + N_{c})/2$.
 For the ESOLV data we used 12 nodes in the input laver corresponding to the 12 input features. eight nodes in the hidden laver. and four in the output laver.," For the ESOLV data we used 12 nodes in the input layer corresponding to the 12 input features, eight nodes in the hidden layer, and four in the output layer."
 For the SDSS data we used five input laver nodes. five hidden laver nodes. aud six output laver nodes.," For the SDSS data we used five input layer nodes, five hidden layer nodes, and six output layer nodes."
"band can be described by a,=0.86+0.07.",band can be described by $\alpha_{o} = 0.86 \pm 0.07$.
 The hieh redshift of lis object is even highly plausible. because it was nof possible to resolve its host galaxy on TST suap shot exposures (Scarpa et al. 1999)).," The high redshift of this object is even highly plausible, because it was not possible to resolve its host galaxy on HST snap shot exposures (Scarpa et al. \cite{scarpa}) )."
 The apparent maeuitude varies slightly through the different «yochs. having reached the faintest value of and in February 1999 (direct oenagineg with Calar Alto 3.511 aud MIOSCA).," The apparent magnitude varies slightly through the different epochs, having reached the faintest value of and in February 1999 (direct imaging with Calar Alto 3.5m and MOSCA)."
 These values were derived by comparison with photometric stand stars in the field of view (Villata et al. 1998))., These values were derived by comparison with photometric standard stars in the field of view (Villata et al. \cite{villata}) ).
 4j=lausectApe+ and yy=0.5 leads to an absolute optical magnitude of at least Mg=27.2mae ae Mps26.1 Gucludiug I-correction).," $H_{0}= 50 \; {\rm km \; sec^{-1} \; Mpc^{-1}}$ and $q_{0} = 0.5$ leads to an absolute optical magnitude of at least $M_{R}= -27.2 \;
{\rm mag}$ and $M_{B} \le -26.4$ (including K-correction)."
 Scarpa οἳ al. (1999)), Scarpa et al. \cite{scarpa}) )
 report the discovery of three arclike structures around 15171656 iu their HIST snapshot survey of DL Lac objects., report the discovery of three arclike structures around 1517+656 in their HST snapshot survey of BL Lac objects.
 The radius of this possible fracimented Eiusteiu ring is 2.1 arcsec., The radius of this possible fragmented Einstein ring is 2.4 arcsec.
 Tf this feature indeed represeuts ai Einstein rius. the mass of tle host ealaxv of 15171656 can casily be estimated.," If this feature indeed represents an Einstein ring, the mass of the host galaxy of 1517+656 can easily be estimated."
 As the redshift of these background objects is not known. we can onlv derive a lower luit for the mass of the lens.," As the redshift of these background objects is not known, we can only derive a lower limit for the mass of the lens."
" For a spherically svuuuetric mass distribution (with 0 being the radius of he Eiustein ring. Dg the augular size distance from the observer to the lens. D, from observer to the source. aud. Dg. the distauce from the leas to the source) we ect (cf."," For a spherically symmetric mass distribution (with $\theta$ being the radius of the Einstein ring, $D_{\rm d}$ the angular size distance from the observer to the lens, $D_{\rm s}$ from observer to the source, and $D_{\rm ds}$ the distance from the lens to the source) we get (cf."
 Schneider et al. 1993)):, Schneider et al. \cite{schneider}) ):
 Thus the lower μπιτ for the mass inside the Eimsteiu ring is A=15-104AL: for Eiustein-de Sitter cosmology and Ty=ακουςtMpe+.," Thus the lower limit for the mass inside the Einstein ring is $M = 1.5 \cdot 10^{12}\,M_{\sun}$ for Einstein-de Sitter cosmology and $H_0=50\,\rm
km\,sec^{-1}\,Mpc^{-1}$."
 For other realistic world models (also inchiding a positive cosmological constant). this luit is even higher.," For other realistic world models (also including a positive cosmological constant), this limit is even higher."
 Asstuuing an isothermal sphere for the lens. the velocity dispersion in the restframe cau be calculated by ludepeudeut of fy we ect ai value of at least 330kinsce+ for Einstein-de Sitter cosmology. and slightly less (320ausec 1) for a flat low-density universe (Q34= 0.3. Oy= 0.7).," Assuming an isothermal sphere for the lens, the velocity dispersion in the restframe can be calculated by Independent of $H_0$ we get a value of at least $330\,\rm km\,sec^{-1}$ for Einstein-de Sitter cosmology, and slightly less $320\,\rm km\,sec^{-1}$ ) for a flat low-density universe $\Omega_{\rm M}=0.3$ , $\Omega_\Lambda=0.7$ )."
 Other models again lead to even higher values., Other models again lead to even higher values.
 The true values of the mass aud velocity dispersion nuelt be much ligher if the redshift of the source is significantly below :z2., The true values of the mass and velocity dispersion might be much higher if the redshift of the source is significantly below $z \approx 2$.
 Figures | and Ὁ show the mass and velocity dispersion as a function of the source redshift., Figures \ref{fig:mass} and \ref{fig:sigma} show the mass and velocity dispersion as a function of the source redshift.
 If the observed absorption is caused wea foreground object and the redshift of 1517 is higher than 0.7. the mass and velocity dispersion of the host galaxy have to be even ligher.," If the observed absorption is caused by a foreground object and the redshift of 1517 is higher than 0.7, the mass and velocity dispersion of the host galaxy have to be even higher."
Zakamskaetal.(2003) sample is strongly biased towards objects with high line luminosities € L[OIIII-1077. erg s+).,\cite{zak03} sample is strongly biased towards objects with high line luminosities ( $>$ $^{42}$ erg $^{-1}$ ).
" L[OIII] ean be as low as —I07"" erg + in radio-quiet type | quasars and narrow line radio galaxies (e.g. Bennert et al. 2002.."," L[OIII] can be as low as $\sim$ $^{40}$ erg $^{-1}$ in radio-quiet type 1 quasars and narrow line radio galaxies (e.g. Bennert et al. \citeyear{ben02},"
 Tadhunter et al. 1998))., Tadhunter et al. \citeyear{tadh98}) ).
 There must be many type 2 quasars with L[OIII] in the range ~ 107?7 erg +., There must be many type 2 quasars with L[OIII] in the range $\sim$ $^{40-42}$ erg $^{-1}$.
 The hybrid objects discussed here. where stellar photoionization is obvious. have among the lowest L[ONT] values within the type 2 quasar sample of Zakamskaetal.(2003).," The hybrid objects discussed here, where stellar photoionization is obvious, have among the lowest L[OIII] values within the type 2 quasar sample of \cite{zak03}."
. This tentatively suggests that stellar photoionization could be relatively more important in type 2 quasars with lower [ONT] luminosities., This tentatively suggests that stellar photoionization could be relatively more important in type 2 quasars with lower [OIII] luminosities.
 If these objects have lower power AGNs (e.g. Simpson 1998)). it would seem plausible to find a relatively higher contribution of the stellar photoionized gas to the observed emission line spectrum.," If these objects have lower power AGNs (e.g. Simpson \citeyear{sim98}) ), it would seem plausible to find a relatively higher contribution of the stellar photoionized gas to the observed emission line spectrum."
 The fraction of hybrid objects. therefore. could be much larger. if lower L[OIII] values were considered.," The fraction of hybrid objects, therefore, could be much larger, if lower L[OIII] values were considered."
 Since star formation has been found in differents classes of type 2 AGNSs at different > (e.g. Holt et al. 2007..," Since star formation has been found in differents classes of type 2 AGNs at different $z$ (e.g. Holt et al. \citeyear{holt07},"
 Tadhunter et al. 2005..," Tadhunter et al. \citeyear{tadh05},"
 Alonso-Herrero et al. 2008) , Alonso-Herrero et al. \citeyear{al08}) )
and in at least a fraction of type 2 quasars nez-Sansigre et al. 2008..," and in at least a fraction of type 2 quasars nez-Sansigre et al. \citeyear{mar08},"
 Lacy et al. 20073).," Lacy et al. \citeyear{lacy07}) ),"
 stellar+AG photoionization is a plausible scenario., stellar+AGN photoionization is a plausible scenario.
 Inspired by these arguments. we have explored an alternative scenario to pure AG photoionization in which a varying contribution of stellar ionized gas is added to the line fluxes in type 2 quasars.," Inspired by these arguments, we have explored an alternative scenario to pure AGN photoionization in which a varying contribution of stellar ionized gas is added to the line fluxes in type 2 quasars."
 Although the hybrid models presented here are rather simplistic and no reliable quantitative results can be extracted about the relative importance of stellar vs. AGN photoionization. they reproduce the type 3 quasar ratios (and. type 2 AGN in general) quite successfuly (better than the standard AGN sequence).," Although the hybrid models presented here are rather simplistic and no reliable quantitative results can be extracted about the relative importance of stellar vs. AGN photoionization, they reproduce the type 2 quasar ratios (and type 2 AGN in general) quite successfully (better than the standard AGN sequence)."
 This suggests that stellar photoionization might also be present in many type 2 quasars. in addition to AGN photoionization.," This suggests that stellar photoionization might also be present in many type 2 quasars, in addition to AGN photoionization."
 Given that other type 2 AGNs have similar line ratios. this applies as well in those cases.," Given that other type 2 AGNs have similar line ratios, this applies as well in those cases."
" On the other hand. and contrary to the more sofisticated AGN models discussed in the literature. the hybrid models cannot solve the “temperature problem"" (see $33.) £9)."," On the other hand, and contrary to the more sofisticated AGN models discussed in the literature, the hybrid models cannot solve the “temperature problem” (see 3.1 $B$ )."
 Regarding all other line ratios. given the strong degeneracy with more sotisticated AGN photoionization models. it is not possible to favour one scenario or another in terms of the line ratios.," Regarding all other line ratios, given the strong degeneracy with more sofisticated AGN photoionization models, it is not possible to favour one scenario or another in terms of the line ratios."
 Other sources of information would be very valuable to test whether stellar photoionization is present in type 2 quasars and characterize its importance relative to the AGN photoionization (e.g. evidence for extended star formation. possible correlations between star formation indicators and the line ratios which require a stronger contribution of stellar ionized gas. lower continuum polarization level in objects with hints of stellar photoionization from the line ratios. ete).," Other sources of information would be very valuable to test whether stellar photoionization is present in type 2 quasars and characterize its importance relative to the AGN photoionization (e.g. evidence for extended star formation, possible correlations between star formation indicators and the line ratios which require a stronger contribution of stellar ionized gas, lower continuum polarization level in objects with hints of stellar photoionization from the line ratios, etc)."
 If the emission lines. in particular [OIII]A3007 have a strong contribution of stellar photoionized gas. the [ΟΠΠ line might not be a good indicator of AGN power.," If the emission lines, in particular $\lambda$ 5007 have a strong contribution of stellar photoionized gas, the [OIII] line might not be a good indicator of AGN power."
 The authors thank to an anonymous referee for useful comments which helped to improve the paper., The authors thank to an anonymous referee for useful comments which helped to improve the paper.
 The work by MVM has been funded with support from the Spanish Ministerio de Educaciónn y Ciencia through the grants AYA2004-02703 and AYA2007-64712. and co-tinanced with FEDER funds.," The work by MVM has been funded with support from the Spanish Ministerio de Educaciónn y Ciencia through the grants AYA2004-02703 and AYA2007-64712, and co-financed with FEDER funds."
 Thanks to Joanna Holt for providing the data set on radio galaxies., Thanks to Joanna Holt for providing the data set on radio galaxies.
 LB was supported by the CONACYT grant J-50296, LB was supported by the CONACyT grant J-50296
"Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with ΝΑΟΙ as a domestic partner, NASA and STFC (UK) as international partners.","	Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, 	NASA and STFC (UK) as international partners."
 Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA., 	Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA.
 team mainly consists of scientists from institutes in the partner countries., 	This team mainly consists of scientists from institutes in the partner countries.
" Support for the post-launch operation is provided by JAXA and ΝΑΟΙ) (Japan). STFC (U.K.), NASA. ESA, and NSC (Norway)."," 	Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA, ESA, and NSC (Norway)."
 We greatly appreciate the proofreading/editing assistance from the GCOE program., 	We greatly appreciate the proofreading/editing assistance from the GCOE program.
nodel produces temperatures which are too low. then real systems will be even more stable against spiral structure growth and (ragmentatiou aud less likely to produce large. coagulated grains.,"model produces temperatures which are too low, then real systems will be even more stable against spiral structure growth and fragmentation and less likely to produce large, coagulated grains."
 If hey are too hieh. the model could inaccurately portray the disk as too stable.," If they are too high, the model could inaccurately portray the disk as too stable."
 Are the temperatures »oduced i1 the nodel too low?, Are the temperatures produced in the model too low?
 We can coustrain the temperatures in the model by coupariug the adiated eilerey from observed systems (here specifically to 5)). to that produced by he simulajon ln vari[9]is waveleneth baucds.," We can constrain the temperatures in the model by comparing the radiated energy from observed systems (here specifically to ), to that produced by the simulation in various wavelength bands."
 This comparison requi‘es that we relate the lumiuous output ) he temperaltre., This comparison requires that we relate the luminous output to the temperature.
 lu 'e0οςous where 1ie optical depth is high. as it is in the ace|jon disks. the radiaed emission can be ap)roximatecl as a blackbody with a temperature of he disk plotosphere.," In regions where the optical depth is high, as it is in the accretion disks, the radiated emission can be approximated as a blackbody with a temperature of the disk photosphere."
 The disks uidplane ane [Rm)10tOsJere teruperatures are then related to each «11er by a given Rosseand opacity aud the loca vertical deusity/temperature profile., The disk's midplane and photosphere temperatures are then related to each other by a given Rosseland opacity and the local vertical density/temperature profile.
 We determite such a p‘olile as by-product. of the cooliug inodel in tlis work. uuder the assuuptiou that the vertical strucure is instantaneously ainbatic.," We determine such a profile as by-product of the cooling model in this work, under the assumption that the vertical structure is instantaneously adiabatic."
 Otjer work (Bellefad.1997:D'Alessioal.L998) las shown that the structure may instead be super aciabatic.," Other work \citep{BCKH,dall98} has shown that the structure may instead be super adiabatic."
 If this is he case hen the midplane temperatu'es will be higher. aud Olr conclusio Lis stroger.," If this is the case then the midplane temperatures will be higher, and our conclusion is stronger."
 The ofyposi)e case may be true itstead: la‘oe relative heating cau occur at high altitu«es. evell hough the μιel altitde heatiug is small in €an absoltte sense (Picketelah2000).," The opposite case may be true instead: large relative heating can occur at high altitudes, even though the high altitude heating is small in an absolute sense \citep{Pick00}."
". This meatus hat vertic:"" eliperaure structure may becje clisto""θε and ai incor‘ect midplane tem»erature could be it[erred.", This means that vertical temperature structure may become distorted and an incorrect midplane temperature could be inferred.
 Hie1 altitude cvuanical ating will play a role simi| 10 hig1 altitude passive eating from stellar potons., High altitude dynamical heating will play a role similar to high altitude passive heating from stellar photons.
 D'Alessioe£al.(1998) show that this procs produces a tem»erature inversion |igh above he photospiere. but his region contributes negliibly to he radiated flux.," \citet{dall98} show that this process produces a temperature inversion high above the photosphere, but this region contributes negligibly to the radiated flux."
 Therelore. we can rel von he modeled racliated outpuU of the simulation to represent accuraterv he temperalires at the clisk miplaue. giveu the physical procOeSSOS duced. i1 the calculatio[u," Therefore, we can rely on the modeled radiated output of the simulation to represent accurately the temperatures at the disk midplane, given the physical processes included in the calculation."
n Our couchsions abou the formation of platets will be confirned if we d that the energy output TOL1 Che disks is equal to οr lesst1hau that observed., Our conclusions about the formation of planets will be confirmed if we find that the energy output from the disks is equal to or less than that observed.
 To oain a valic CODDparlsol between the oserved aiCL nucxleled {1wes. we Dnuust be certain 1le Ix. [roin οἱher parts of the system (e.g. e οἱuary disX aud envelope) is not a iLicfall contributo “to the observed flux. used i e comparison.," To obtain a valid comparison between the observed and modeled fluxes, we must be certain that the flux from other parts of the system (e.g. the circumbinary disk and envelope) is not a significant contributor to the observed flux used in the comparison."
" We 1iist also be certain that 1101 between the source aud the obse""ver not altered the e1nted flux.", We must also be certain that extinction between the source and the observer has not altered the emitted flux.
 We therefore ue very liigh spaial resolution photoiet Iva £& wavelenehs. not alected by extinction.," We therefore require very high spatial resolution photometry at long wavelengths, not affected by extinction."
 h igure κ)+)3.. the|ughest available resoluion. lonees waveleuegtl observatious of aare poOtted aud comyared to the flux. densities prodced from tlje siniulatjon.," In figure \ref{fig:freq-cmp}, the highest available resolution, long wavelength observations of are plotted and compared to the flux densities produced from the simulation."
 For all wavelenetls betwee lHN»!+) cin aud 870. jan. the observed {wes exceed those oplained from the simulation by a factor of 5.," For all wavelengths between 1.3 cm and 870 $\mu$ m, the observed fluxes exceed those obtained from the simulation by a factor of $\sim 5$."
 The differences at 1.1 nuu aud STO pumare larger. a factor «10. however they co not resolve the binary aud may contain some co(αμαug [lux frou the cireumbiuary e1vironnent.," The differences at 1.4 mm and 870 $\mu$ m are larger, a factor $\sim10$, however they do not resolve the binary and may contain some contaminating flux from the circumbinary environment."
 Civen tlese comparisons. we can conclude hat the temperatures iu the disk are coiservatively estimated by the simulation (too low) aud tlat the disks in aare no'e stable against spiral ari growth aud fragmentation aud less likely to allow cus coagulation," Given these comparisons, we can conclude that the temperatures in the disk are conservatively estimated by the simulation (too low) and that the disks in are more stable against spiral arm growth and fragmentation and less likely to allow dust coagulation"
"following periastron contrasts with the simulations of shorter period OB star systems presented by (2009),, in which clumps formed close to periastron survived until apastron.","following periastron contrasts with the simulations of shorter period OB star systems presented by , in which clumps formed close to periastron survived until apastron."
" The growth of NTSIs corrugates the WCR apex, increasing the shock obliquity, thereby reducing the efficiency with which wind kinetic energy is thermalized."," The growth of NTSIs corrugates the WCR apex, increasing the shock obliquity, thereby reducing the efficiency with which wind kinetic energy is thermalized."
" Consequently, postshock gas temperatures are further reduced and cooling becomes even more effective."," Consequently, postshock gas temperatures are further reduced and cooling becomes even more effective."
 The runaway disruption continues as periastron is approached and the WCR apex oscillates wildly in the vicinity of the O star until a collision occurs at @~0.95., The runaway disruption continues as periastron is approached and the WCR apex oscillates wildly in the vicinity of the O star until a collision occurs at $\phi \simeq 0.95$.
" The preshock O star’s wind, which suffers substantial radiative inhibition at phases close to periastron, has little room to accelerate and its ram pressure is unable to hold back the incoming WR wind, leading to a of the WCR onto the O star at ¢=0.96."," The preshock O star's wind, which suffers substantial radiative inhibition at phases close to periastron, has little room to accelerate and its ram pressure is unable to hold back the incoming WR wind, leading to a of the WCR onto the O star at $\phi=0.96$."
 We note that contrary to the predictions of Fig., We note that contrary to the predictions of Fig.
 3 radiative braking does not prevent a collapse of the WCR because the region over which braking is predicted to occur is occupied by postshock gas at T>10° K which we assume to be mostly ionized and consequently to be subject to a negligible driving force., \ref{fig:radiative_braking} radiative braking does not prevent a collapse of the WCR because the region over which braking is predicted to occur is occupied by postshock gas at $T>10^{6\;}$ K which we assume to be mostly ionized and consequently to be subject to a negligible driving force.
" However, the strength of radiative braking is dependent on the parameters adopted, i.e. the k and α1997).."," However, the strength of radiative braking is dependent on the parameters adopted, i.e. the $k$ and $\alpha$."
 When calculating the line force applied to the WR’s wind by the O star’s radiation we use the O star’s parameters (ko=0.3 and ao= 0.52).," When calculating the line force applied to the WR's wind by the O star's radiation we use the O star's parameters $k_{\rm O}=0.3$ and $\alpha_{\rm
 O}=0.52$ )."
" Hence, if the WR's radiation-wind coupling was adopted (kwr=0.42 and awr= 0.47) the decelerative line force would be greater and radiative braking may be more effective than in our calculations."," Hence, if the WR's radiation-wind coupling was adopted $k_{\rm WR}=0.42$ and $\alpha_{\rm WR}=0.47$ ) the decelerative line force would be greater and radiative braking may be more effective than in our calculations."
" Interestingly, as the stars reach ¢=1.0 the apex of the WCR is essentially stabilised against NTSIs by the continuing collapse (i.e. the ram pressure of the WR's wind pins the WCR apex to the O star preventing oscillations) and only relatively small amplitude KH instabilities can be seen in the trailing arm of the WCR (Fig. 6))."," Interestingly, as the stars reach $\phi=1.0$ the apex of the WCR is essentially stabilised against NTSIs by the continuing collapse (i.e. the ram pressure of the WR's wind pins the WCR apex to the O star preventing oscillations) and only relatively small amplitude KH instabilities can be seen in the trailing arm of the WCR (Fig. \ref{fig:collapse_images}) )."
" Despite the O star's wind being overwhelmed between the stars it continues to drive a wind on the side facing the WR and, as in model A, carves a tenuous cavity through the dense WR wind (Fig. 5))."," Despite the O star's wind being overwhelmed between the stars it continues to drive a wind on the side facing the WR and, as in model A, carves a tenuous cavity through the dense WR wind (Fig. \ref{fig:driven_images}) )."
" As the wind driven from the far side of the O star shocks against the leading arm of the WCR (which becomes wrapped around the stars) it reheats the postshock gas, and consequently the leading arm of the WCR is hotter than the trailing arm (see the temperature plots at 9=1.0 and 1.1 in Figs."," As the wind driven from the far side of the O star shocks against the leading arm of the WCR (which becomes wrapped around the stars) it reheats the postshock gas, and consequently the leading arm of the WCR is hotter than the trailing arm (see the temperature plots at $\phi=1.0\;$ and 1.1 in Figs."
 4 and 5))., \ref{fig:vterm_images} and \ref{fig:driven_images}) ).
" Following periastron, as the separation of the stars expands, the preshock ram pressure of the WR's wind - which pins the shocks onto the O star - decreases."," Following periastron, as the separation of the stars expands, the preshock ram pressure of the WR's wind - which pins the shocks onto the O star - decreases."
" Therefore, when it weakens sufficiently, the WCR is permitted to move away from the O star."," Therefore, when it weakens sufficiently, the WCR is permitted to move away from the O star."
" Simultaneously, as the stars recede, radiative inhibition decreases and the acceleration of the O star's wind increases."," Simultaneously, as the stars recede, radiative inhibition decreases and the acceleration of the O star's wind increases."
statistical equilibrium caleulations for some molecular lines (rotaetal.2002.2004:Ii-rola&Yamamoto2006:Aikawaetal. 2005).,"statistical equilibrium calculations for some molecular lines \citep{hirota2002, hirota2004, hirota2006, aikawa2005}."
. Although we cannot directly compare the ll» densiües in CD130-3 and L673-5MMMA. intense spectra of hieh-censitv (racers such as HCO imply the Ho densities of an order of 10? +.," Although we cannot directly compare the $_{2}$ densities in CB130-3 and L673-SMM4, intense spectra of high-density tracers such as $^{13}$ $^{+}$ imply the $_{2}$ densities of an order of $>$ $^{5}$ $^{-1}$."
 Therefore. the CCPRs themselves are not just lower extinction and less dense eas clumps but rather chemically vounger dense cores compared wilh other prestellar cores such as L1544 and L1498.," Therefore, the CCPRs themselves are not just lower extinction and less dense gas clumps but rather chemically younger dense cores compared with other prestellar cores such as L1544 and L1498."
 Further observations of (hese regions including less-dense peripheries around dense cores will be necessary to explore a role of a surrounding environment in formation of dense cores including CCPRs., Further observations of these regions including less-dense peripheries around dense cores will be necessary to explore a role of a surrounding environment in formation of dense cores including CCPRs.
 Existence of CCPRs outside the Taurus region indicates that the origin of CCPRs is nol ascribed to the regional specialty of (he Taurus cloud., Existence of CCPRs outside the Taurus region indicates that the origin of CCPRs is not ascribed to the regional specialty of the Taurus cloud.
 Hence. our results may strengthen ihe argument that the evolutionary timescale of dense cores is different [rom cloud to cloud. as proposed bv Iirotaetal.(2009) ancl Sakaietal.(2009).," Hence, our results may strengthen the argument that the evolutionary timescale of dense cores is different from cloud to cloud, as proposed by \citet{hirota2009} and \citet{sakai2009}."
. We recently. found. another ivpeof carbon-chain rich core in the Lupus Molecular. Cloud named. Lupus-1AÀ (Sakai2010:Shiinoetal. 2011).," We recently found another typeof carbon-chain rich core in the Lupus Molecular Cloud named Lupus-1A \citep{sakai2010, shiino2011}."
. This source shows extraordinarilv intense spectra of Cyl and longer carbon-chain molecules. while that of CCS is not as bright as in the known CCDPIs including CD130-3 and L67T3-SMMA.," This source shows extraordinarily intense spectra of $_{4}$ H and longer carbon-chain molecules, while that of CCS is not as bright as in the known CCPRs including CB130-3 and L673-SMM4."
. All of these sources will be good targets to reveal chemical differentiation among the large-scale molecular cloud complexes., All of these sources will be good targets to reveal chemical differentiation among the large-scale molecular cloud complexes.
 In addition. they can be useful and unique sources to investigate formation and evolution of carbou-chain molecules in dark cloud cores in the chemically voung evolutionary phase.," In addition, they can be useful and unique sources to investigate formation and evolution of carbon-chain molecules in dark cloud cores in the chemically young evolutionary phase."
 Thev will shed light on the initial state of chemical and. dyiamical evolution of dark eloud cores by. the high-resolution observations with Atacama Large Millimeter/Submillimeter Arrav (ALMA)., They will shed light on the initial state of chemical and dynamical evolution of dark cloud cores by the high-resolution observations with Atacama Large Millimeter/Submillimeter Array (ALMA).
 The 45 m radio telescope is operated by Nobevama adio Observatory. a branch of National Astronomical Observatory of Japan (NÀOJ).," The 45 m radio telescope is operated by Nobeyama Radio Observatory, a branch of National Astronomical Observatory of Japan (NAOJ)."
 We are grateful to all the stall of Nobevama Raclio Observatory of NAOJ ancl Effelsberg Observatory of MPIIR. for their assistance in observations., We are grateful to all the staff of Nobeyama Radio Observatory of NAOJ and Effelsberg Observatory of MPIfR for their assistance in observations.
 TI] and NS thank to the Inoue Foundation lor Science for the financial support (Research Aid of Inoue Foundation [for Science)., TH and NS thank to the Inoue Foundation for Science for the financial support (Research Aid of Inoue Foundation for Science).
 This study is partly supported by Grant-in-Aid from The Ministry of Education. Culture. Sports. Science and Technology of Japan (No.," This study is partly supported by Grant-in-Aid from The Ministry of Education, Culture, Sports, Science and Technology of Japan (No."
 21224002 and 21740132). No:dom.Ellelsbere.., 21224002 and 21740132). .
distinguished as the dominaut outflow mechanisin.,distinguished as the dominant outflow mechanism.
 We extend our thanks to Carlo Staughellini. Istituto di Raclioastronomia del C.N.R. Bologna. Italy. for supplying the 5 GIIz radio data for PISS 1315112. and to Tasso Tziounis. CSIRO Raciophysics Laboratory. Pembroke and Vimiera Ros. Miursfield. NSW. Australia. for retrieving the radio data for PINS 1519-79 from Edward iuss thesis.," We extend our thanks to Carlo Stanghellini, Istituto di Radioastronomia del C.N.R, Bologna, Italy, for supplying the 5 GHz radio data for PKS 1345+12, and to Tasso Tzioumis, CSIRO Radiophysics Laboratory, Pembroke and Vimiera Rds, Marsfield, NSW, Australia, for retrieving the radio data for PKS 1549-79 from Edward King's thesis."
 We also thank the anonviuous referee for their useful commucuts., We also thank the anonymous referee for their useful comments.
 Support for Proposal umber HST-CO-09LOLLOA was provided bv NASA through a graut from the Space Telescope Science. Tustitute. which is operated by the Association of Universities for Researcli in Astrouoniv. Iucorporated. under NASA coutract NAS5-26555.," Support for Proposal number HST-GO-09401.10A was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
in studying binary stars. and have long been argued to be well-suited to studying extra- planets (Colavita&Shao1994:EisnerKulkarni2001).. to date results using these techniques have been limited (Benedictetal.2002b).,"in studying binary stars, and have long been argued to be well-suited to studying extra-solar planets \citep{col94b,eis01}, to date results using these techniques have been limited \citep{ben02b}."
. There are several reasons why it is desirable to develop viable astrometric methods., There are several reasons why it is desirable to develop viable astrometric planet-detection methods.
 Most importantly. the parameter space explored by astrometry is complementary to (hat of radial velocity (astrometry is more sensitive to larger separations).," Most importantly, the parameter space explored by astrometry is complementary to that of radial velocity (astrometry is more sensitive to larger separations)."
 Second. unlike current racial velocity detections. astrometric techniques can be used to determine (he orbital inclination of a planet.," Second, unlike current radial velocity detections, astrometric techniques can be used to determine the orbital inclination of a planet."
 Finally. astrometry is parüceularlw well-suited to studying binary stellar svstems: such svstenis challenge other planet-finding techniques.," Finally, astrometry is particularly well-suited to studying binary stellar systems; such systems challenge other planet-finding techniques."
 For example. radial velocimetry can suffer [rom systematic velocity errors caused by spectral contamination from the light of the second star (Vogtetal.2000)..," For example, radial velocimetry can suffer from systematic velocity errors caused by spectral contamination from the light of the second star \citep{Vog:00::}."
 Similar problems are [aced by coronographic techniques. where the light from the second star is not usually blocked by the occulting mask.," Similar problems are faced by coronographic techniques, where the light from the second star is not usually blocked by the occulting mask."
 In this paper we describe recent efforts to obtain very high precision narrow-angle astromelrv using PTI (o observe binary stars wilh separations less than one arcsecond. i.e. svslenms (hat are (vpically observed. using speckle interferometry (Sala2002) or adaptive oplies.," In this paper we describe recent efforts to obtain very high precision narrow-angle astrometry using PTI to observe binary stars with separations less than one arcsecond, i.e. systems that are typically observed using speckle interferometry \citep{sah02} or adaptive optics."
 Such small separations allow us (o achieve astrometric precision on (he order of LO pas. Which for a typical binary svstem in our target sample (binary separation of 20 AU). should allow us to detect planets with masses down to 0.5 Jupiter masses in orbils in the 2 AU range.," Such small separations allow us to achieve astrometric precision on the order of 10 $\mu$ as, which for a typical binary system in our target sample (binary separation of 20 AU), should allow us to detect planets with masses down to 0.5 Jupiter masses in orbits in the 2 AU range."
 This approach has been suggested. CIraub.Carleton&Porro1996). and tried (Dvck.Benson&Schloerb1995:Dagnuoloetal.2003) before. though with limited precision.," This approach has been suggested \citep{tcp96} and tried \citep{dyck95,chara03} before, though with limited precision."
 llowever. this work is unique in that it makes use of a phase-tracking interferometer: the use of phase-referencing (Lane&Colavila2003) removes much of the effect of atmospheric turbulence. improving the astrometric precision by a [actor of order 100.," However, this work is unique in that it makes use of a phase-tracking interferometer; the use of phase-referencing \citep{lc03} removes much of the effect of atmospheric turbulence, improving the astrometric precision by a factor of order 100."
 The Palomar Testbed Interferometer (PTI) is located on Palomar Mountain near San Diego. CA (Colavitaetal.1999).," The Palomar Testbed Interferometer (PTI) is located on Palomar Mountain near San Diego, CA \citep{col99}."
. It was developed by the Jet Propulsion Laboratory. California. Institute of Technology for NASA. as a testbed for interlerometric techniques applicable to the Neck Interferometer aud other missions such as the Space Interferometry Mission. SIM.," It was developed by the Jet Propulsion Laboratory, California Institute of Technology for NASA, as a testbed for interferometric techniques applicable to the Keck Interferometer and other missions such as the Space Interferometry Mission, SIM."
 It operates in the J 1 μαι). (1 μι) and Ix. (2.2;0n) bands. and combines starlight [rom two out of three available 40-cm apertures.," It operates in the J $1.2 \mu{\rm
m}$ ),H $1.6 \mu{\rm m}$ ) and K $2.2 \mu{\rm m}$ ) bands, and combines starlight from two out of three available 40-cm apertures."
 The apertures form a (rianele with 86 and 110 meter baselines., The apertures form a triangle with 86 and 110 meter baselines.
 The paper is organized as follows: in Section 2 we describe (he experiment and derive expected perlormance levels., The paper is organized as follows: in Section 2 we describe the experiment and derive expected performance levels.
 In Section 3 we describe initial observations as well as the extensive data analvsis processing required (o achieve the desired. astrometric precision., In Section 3 we describe initial observations as well as the extensive data analysis processing required to achieve the desired astrometric precision.
 In section 4 we discuss our preliminary results. ancl in Section 5 we discuss the prospects of a larger search.," In Section 4 we discuss our preliminary results, and in Section 5 we discuss the prospects of a larger search."
As shown in Table 1 NO-3 b has the largest predicted of all currently known cxoplanets with ο>0.1.,As shown in Table \ref{tab:prec} XO-3 b has the largest predicted of all currently known exoplanets with $e>0.1$.
 While its eccentricity is uot as high as that of ILAT-P-2 b. it is more massive and orbits closer to its host star NO-3 (also shown as GSC 03727-01061).," While its eccentricity is not as high as that of HAT-P-2 b, it is more massive and orbits closer to its host star XO-3 (also known as GSC 03727-01064)."
" Astuning he,=0.25. the expected value of iua d8 ~11 deg/ceutiry. about three times as muuch as the contribution from GR."," Asuming $k_{2,p}=0.25$, the expected value of $\dot{\omega}_{\rm tide}$ is $\sim 11$ deg/century, about three times as much as the contribution from GR."
 Tu order to assess the detectability of i with radial velocities for NO-3 b we need to know oj. but uufortuuatelv he precision of the racial velocity micasurements presented in? is too coarse to allow a determination of this quautity (To.2LOO i 1j.," In order to assess the detectability of $\dot{\omega}$ with radial velocities for XO-3 b we need to know $\sigma_{\rm jitter}$, but unfortunately the precision of the radial velocity measurements presented in \citet{Johns-Krull2008a} is too coarse to allow a determination of this quantity $\sigma_{\rm obs} \gtrsim 100$ m $^{-1}$ )."
 Based on the spectral type ΕΟΝ and sins=18.540.2 kin bof NO-3 b (2) we can expect it to have a rather high value of stellar jitter Gitta230 ni sec1o).," Based on the spectral type F5V and $v\sin i=18.5\pm 0.2$ km $^{-1}$ of XO-3 b \citep{Johns-Krull2008a} we can expect it to have a rather high value of stellar jitter $\sigma_{\rm jitter}
\gtrsim 30$ m $^{-1}$ \citep{Saar1998a}."
 Therefore. aud just as is the case fox ILAT-P-2 » we do not expect w to be detectable with radial velocity observatious due to the expected jitter.," Therefore, and just as is the case for HAT-P-2 b, we do not expect $\dot{\omega}$ to be detectable with radial velocity observations due to the expected jitter."
 The value of uw derived for NO-3 b is cousistent with 0 and we therefore would not expect to see variations in tle ine between primary and secondary trausits in case the secondary. transit was observed (see Figure 1))., The value of $\omega$ derived for XO-3 b is consistent with 0 and we therefore would not expect to see variations in the time between primary and secondary transits in case the secondary transit was observed (see Figure \ref{fig:dDt}) ).
 As wz0. the rueanomaly at the time of trausit will be f;z7/2 or 37/2.," As $\omega \approx 0$, the trueanomaly at the time of transit will be $f_t \approx \pi/2$ or $3\pi/2$."
 Given that the impact parameter is bzz(0.8 we ect from Equation 3.2 that [dnD/dt|~3.56107 ceutury|.," Given that the impact parameter is $b\approx 0.8$ we get from Equation \ref{eq:dlnD} that $|d\ln D/dt|
\sim 3.56\times10^{-2}$ $^{-1}$."
 Using the fact that D=0.11 days we expect then a change of ~[3 κος in D over 10 vears., Using the fact that $D = 0.14$ days we expect then a change of $\sim 43$ sec in $D$ over 10 years.
 Just as is the case for ILAT-P-2 b. this could be detectable with determinations of D to within a few seconds and there is no suitable first epoch vet in haud.," Just as is the case for HAT-P-2 b, this could be detectable with determinations of $D$ to within a few seconds and there is no suitable first epoch yet in hand."
 Iu this work we have studied the observabilitv ofthe precession ofperiastra caused by eeucral relativity in exoplaucts., In this work we have studied the observability of the precession of periastra caused by general relativity in exoplanets.
 We additionally consider the precession caused by tidal deformations and planetary perturbers. which cau produce a precession of comparable or greater magnitude.," We additionally consider the precession caused by tidal deformations and planetary perturbers, which can produce a precession of comparable or greater magnitude."
 We consider radial velocities aud transit ght curve observations aud conclude that for some methods precessious of the magnitude expected from GR will be detectable in timescales of LO vears or less for some close-in. eccentric svstenis.," We consider radial velocities and transit light curve observations and conclude that for some methods precessions of the magnitude expected from GR will be detectable in timescales of $\sim 10$ years or less for some close-in, eccentric systems."
 Iu iore detail. we find that: Iu order to contrast anv detected change in the transit duration rofsubsec:D)) or the time between primary aud secondary refsubsec:Deltat)) to the predictions of given mecliauisin oue needs to know the ecceutricity and longitude of periastrou of the systeis. for which radial velocitiesa are uceded (although uot necessarily of the precision required to directly detect changes in iw with them!?)).," In more detail, we find that: In order to contrast any detected change in the transit duration \\ref{subsec:D}) ) or the time between primary and secondary \\ref{subsec:Deltat}) ) to the predictions of a given mechanism one needs to know the eccentricity and longitude of periastron of the systems, for which radial velocities are needed (although not necessarily of the precision required to directly detect changes in $\omega$ with )."
 Conversely. photometric monitoring of primary trausits are useful iu order to elucidate the nature of a detected change in w by probing for the presence of transit fine variations.," Conversely, photometric monitoring of primary transits are useful in order to elucidate the nature of a detected change in $\omega$ by probing for the presence of transit time variations."
 The presence of the latter would miplv that at least part of any observed changes iu iw could have been produced by additional planetary counrpanions refsecidistinguish))., The presence of the latter would imply that at least part of any observed changes in $\omega$ could have been produced by additional planetary companions \\ref{sec:distinguish}) ).
 Precession of periastra caused by planetary perturbers aud the effects of tidal deformations can be of comparable magnitude to that caused by GR heef , Precession of periastra caused by planetary perturbers and the effects of tidal deformations can be of comparable magnitude to that caused by GR \\ref{ssec:add_prec}) ).
"feetso ftidaldefor niationsoutheprecessionof periastrainparticulurinagbeofthesaimcemagnitudcordoimin in the regine where rofsseciadd,rec)).Tthe CR effects are detectable rofssecxel,og)).", The effects of tidal deformations on the precession of periastra in particular may be of the same magnitude or dominate the total $\dot{\omega}$ in the regime where the GR effects are detectable \\ref{ssec:rel_mag}) ).
"M bilethislinitstheabilitytodirectlycetracttheprocessionductoG Bgiventheuncertaintginthecvpectedpreccssion, "," While this limits the ability to directly extract the precession due to GR given the uncertainty in the expected precession from tides, it might allow to study the tidally induced precession by considering the residual precession after subtracting the effects of GR."
The upcoming, The latter possibility is particularly attractive in systems where the tidally induced precession may dominate the signal.
Aepler aission e," We note that even without considering the confusing effects of tidal contributions to the precession, a measurement of as described in this work would not be competitive in terms of precision with binary pulsar studies \citep[see, e.g.,][for a review ]{Will2006a} and would therefore not offer new tests of GR."
xpects to find a large number of massive plaucts transiting close to their lost stars (?2).. some of which will certainly have significant ecceutricitics.," The upcoming mission expects to find a large number of massive planets transiting close to their host stars \citep{Borucki2003a}, some of which will certainly have significant eccentricities."
 Furthermore. svstenis observed bvAvepler will be extensively monitored for variations in thei transiting time periods in order to search for terrestrialinass plancts using trausit-time variations.," Furthermore, systems observed by will be extensively monitored for variations in their transiting time periods in order to search for terrestrial-mass planets using transit-time variations."
 We have shown that modeling of the transit time durations and further characterization of close-in. eccentric systems miüeht need to take iuto account the effects of COR and tidal deformations as they will become detectable on timescales comparable to the Lvear lifetiiie of the mission. aud certainly on follow-up studies after the mission euds.," We have shown that modeling of the transit time durations and further characterization of close-in, eccentric systems might need to take into account the effects of GR and tidal deformations as they will become detectable on timescales comparable to the 4-year lifetime of the mission, and certainly on follow-up studies after the mission ends."
 We have also shown that planetary companions with super-Earth masses may be detectable by Kepler by the change in transit durations thev induce L2))., We have also shown that planetary companions with super-Earth masses may be detectable by by the change in transit durations they induce \ref{ssec:tpl}) ).
 Additionally. well sampled radial velocity curves," Additionally, well sampled radial velocity curves"
llere we show that gas drag in a laminar disk causes global redistribution and concentration of small solids as they. inspiral.,Here we show that gas drag in a laminar disk causes global redistribution and concentration of small solids as they inspiral.
 It is surprising that this straightforwarc and robust ellect has not been studied before. but the more complex case of the global evolution of solids in turbulent disks has been studied in numerical simulations (Stepinski&Valageas1996).," It is surprising that this straightforward and robust effect has not been studied before, but the more complex case of the global evolution of solids in turbulent disks has been studied in numerical simulations \citep{sv96}."
. In our case. the concentration of nim-sized solids occurs on ~LO? vear timescales. implving that chondrules could play a crucial role in triggering planetesimal formation and in the observed disappearance of dust disks around ‘T Tauri stars.," In our case, the concentration of mm-sized solids occurs on $\sim 10^6$ year timescales, implying that chondrules could play a crucial role in triggering planetesimal formation and in the observed disappearance of dust disks around T Tauri stars."
 For particles with a radius a«9A/4. where A is the gas mean free path. eas drag follows Epstein's law (Weidenschilling1977).," For particles with a radius $a < 9 \lambda/4$, where $\lambda$ is the gas mean free path, gas drag follows Epstein's law \citep{wei77}."
. For the MSN. 9A/4eGt!em. so particles up to chondrule sizes can salelv be treatecl with Epstein drag for all but the innermost regions.," For the MSN, $9 \lambda/4 \simeq .6 \varpi^{11/4}\cm$, so particles up to chondrule sizes can safely be treated with Epstein drag for all but the innermost regions."
 The drift speed due to Epstein drag in a lvelrostatic gas disk is: ∖∖⇁↥∐↲↕⋅≼↲∕⋝∖↿≡∕↗⋝∖∩∕∕∕⋖↜∕↗≜≟∣↽⇀≜≟↕⋟≪⊥∕∕∕≤≥↕⋟∖⊽⊔∐↲⋟∖⊽↥∪↕↽≻↕↽≻↕∐≸≟∐∐∐↲⋅↴∏∐↲∐∏∐∐↲↕," The drift speed due to Epstein drag in a hydrostatic gas disk is: where $t_{\rm st} \equiv \rho_{\rm s}a/(\rho_{\rm g}c_{\rm g}) \ll
1/\Omega$ is the stopping time."
"⋅↕≺∢≀↕↴↥∖↽≀↧↴↥⋯↲⋟∖⊽↕∐⋖⋡∃⋅−↕⋝⇄⋝≀↕↴↕↽≻↕↽≻↥⋡∖↽ to the MSN: more generally the radial dependence goes as eq,x7 where: where we again specilv p as the surface density powerlaw of the gas disk: X,xrf.", The numerical values in \ref{vdr}) ) apply to the MSN; more generally the radial dependence goes as $v_{\rm dr} \propto r^d$ where: where we again specify $p$ as the surface density powerlaw of the gas disk: $\Sigma_{\rm g} \propto r^{-p}$.
 Note that d depends only on gas properties. which we assume to be time constant. and nol on the evolving surface clensity of the solids.," Note that $d$ depends only on gas properties, which we assume to be time constant, and not on the evolving surface density of the solids."
" We will set the inerial factor p,/pzz1 in (22)) because this approximation simplifies ihe mathematics al no sienilicant cost in realism.", We will set the inertial factor $\rho_{\rm g}/\rho \approx 1$ in \ref{vdr}) ) because this approximation simplifies the mathematics at no significant cost in realism.
 Since GI would occur (by the saturation mechanism of re[secisat)) M pa>py. (he procedure represents a factor e2 error al worst in non-critical circumstances.," Since GI would occur (by the saturation mechanism of \\ref{sec:sat}) ) if $\rho_{\rm d} > \rho_{\rm g}$, the procedure represents a factor $\sim 2$ error at worst in non-critical circumstances."
 We assume (hat anv variation in drill speeds associated wilh variations in Pg/ pis hidden by the larger spread that occurs when we have a spectrum of particle sizes., We assume that any variation in drift speeds associated with variations in $\rho_g/\rho$ is hidden by the larger spread that occurs when we have a spectrum of particle sizes.
" With this simplifvine assumption. cg, depends on the solid density and size of the particles. but not on their surface or space density."," With this simplifying assumption, $v_{\rm dr}$ depends on the solid density and size of the particles, but not on their surface or space density."
 An axisvmametrie distribution of uniformly sized particles with surlace clensity Xr./) ," An axisymmetric distribution of uniformly sized particles with surface density $\Sigma(r,t)$ "
accurately cophase our telescope array in order to enhance the dynamics of our instrument.,accurately cophase our telescope array in order to enhance the dynamics of our instrument.
 This study has been financially supported by. CNIZS. INSU and Thales Alenia Space in the frame of different contracts.," This study has been financially supported by CNES, INSU and Thales Alenia Space in the frame of different contracts."
 Our thanks go to Enmmanuclle Abbott for. her help in writing this paper and Alain Dexet for the fabrication of the mechanical parts of our experiment., Our thanks go to Emmanuelle Abbott for her help in writing this paper and Alain Dexet for the fabrication of the mechanical parts of our experiment.
"(2005) have suggested that yields from massive stars are the overwhelming source of carbon at early stages, while later on, at the end of their slower evolution, low- and intermediate-mass stars would be able to contribute carbon ejecta into the ISM in a comparable amount.","(2005) have suggested that yields from massive stars are the overwhelming source of carbon at early stages, while later on, at the end of their slower evolution, low- and intermediate-mass stars would be able to contribute carbon ejecta into the ISM in a comparable amount."
" In their models, a combination of sources differing in mass and thus contributing at different times is required to match the observed trends."," In their models, a combination of sources differing in mass and thus contributing at different times is required to match the observed trends."
" In contrast, Gavilánn, Buell Mollá (2005) have argued that low- and intermediate-mass stars alone may account for the carbon evolution."," In contrast, Gavilánn, Buell Mollá (2005) have argued that low- and intermediate-mass stars alone may account for the carbon evolution."
" Given the very low metal-content of most stars in our sample, with oxygen abundances as low as 10-7? of the solar value, these objects are presumably associated with very early star formation in our Galaxy, before the end of the halo build-up."," Given the very low metal-content of most stars in our sample, with oxygen abundances as low as $10^{-3}$ of the solar value, these objects are presumably associated with very early star formation in our Galaxy, before the end of the halo build-up."
" Thus, by being the only survivors to the present, they provide a window on nucleosynthetic processes taking place in the stars that have existed at such early times."," Thus, by being the only survivors to the present, they provide a window on nucleosynthetic processes taking place in the stars that have existed at such early times."
" A straightforward interpretation of the high C/O values we find at low metallicities is that the first episodes of star formation in the Galaxy provided a source of high C abundance, perhaps thanks to a primordial generation of massive, zero-metallicity stars."," A straightforward interpretation of the high C/O values we find at low metallicities is that the first episodes of star formation in the Galaxy provided a source of high C abundance, perhaps thanks to a primordial generation of massive, zero-metallicity stars."
" According to current theoretical models describing the yields of these hypothetical objects, it is plausible that they could have indeed contributed large C yields (e.g. Chieffi Limongi 2002, 2004)."," According to current theoretical models describing the yields of these hypothetical objects, it is plausible that they could have indeed contributed large C yields (e.g. Chieffi Limongi 2002, 2004)."
 Akerman et al. (, Akerman et al. (
2004) constructed GCE models in order to interpret their tentative discovery of a [C/O] rise at low metallicity.,2004) constructed GCE models in order to interpret their tentative discovery of a [C/O] rise at low metallicity.
" By adopting the Population III yields of Chieffi Limongi (2002), they could reproduce the observed behaviour, in particular when using a top-heavy IMF."," By adopting the Population III yields of Chieffi Limongi (2002), they could reproduce the observed behaviour, in particular when using a top-heavy IMF."
" This would imply, as assumed in the derivation of those yields, that the nucleosynthetic channel 1?C(a,4)190 proceeds at a lower rate in such primordial objects."," This would imply, as assumed in the derivation of those yields, that the nucleosynthetic channel $^{12}$ $\alpha$ $\gamma$ $^{16}$ O proceeds at a lower rate in such primordial objects."
Tn recent vears. a large uunber of disks characterized by a lack of significant mid-intrared (IR) euission aud a rise oeito the far-IR have been detected(e.g.2?)..,"In recent years, a large number of disks characterized by a lack of significant mid-infrared (IR) emission and a rise into the far-IR have been detected\citep[e.g.][]{Brown2007,Merin2010}."
 These are the so-calle ‘transitional disks. and they are thought to be in an intermediate evolutionary state between primordial Class ID protoplanctary disks aud Class III debris disks.," These are the so-called `transitional disks', and they are thought to be in an intermediate evolutionary state between primordial Class II protoplanetary disks and Class III debris disks."
" The lack. of iid-IR excess du cold disks has heen interpreted as a sien of dust cleariug, which can result lu gaps or roles within the disk."," The lack of mid-IR excess in cold disks has been interpreted as a sign of dust clearing, which can result in gaps or holes within the disk."
 These gaps aud holes can be created by several mechanisms. such as a close stellar companion. disk photoevaporatiou. erain growth or v planet formed within the disk.," These gaps and holes can be created by several mechanisms, such as a close stellar companion, disk photoevaporation, grain growth or a planet formed within the disk."
 A plauet formine within the disk is expected to generate a eap while he dust aud eas is accreted onto its surface. sweeping out the orbital region (6.8.7)..," A planet forming within the disk is expected to generate a gap while the dust and gas is accreted onto its surface, sweeping out the orbital region \citep[e.g.][]{Lubow1999}."
 Iu this work. we present high angular resolution deep IR observations of T Cha. a young star with a cold disk.," In this work, we present high angular resolution deep IR observations of T Cha, a young star with a cold disk."
 Its spectral energy. distribution (SED) shows a simall IR excess between 104421 aud a very steep rise between 30;uu. The SED has only been successfully modeled by includiug a gap from 0.2 to AU (2?)..," Its spectral energy distribution (SED) shows a small IR excess between $\mu$ m and a very steep rise between $\mu$ m. The SED has only been successfully modeled by including a gap from 0.2 to AU \citep{Brown2007,Schisano2009}."
 Tu fact. an iuner dusty disk las recently been detected by ?..," In fact, an inner dusty disk has recently been detected by \citet[][]{Olofsson2011}."
 Because one of the possibilities is that the eap has been cleared by a very low-1nass object. we obtained adaptive optics (AO) sparse aperture masking (SAM) observations of T Cha aimed at detecting faint companions within the disk gap.," Because one of the possibilities is that the gap has been cleared by a very low-mass object, we obtained adaptive optics (AO) sparse aperture masking (SAM) observations of T Cha aimed at detecting faint companions within the disk gap."
 T Cha is a high probability member of the voune € Cha association (?).., T Cha is a high probability member of the young $\epsilon$ Cha association \citep{Torres2008}.
" It is a G8-tvpo star with a mass of ~ MM... classified. as a woeak-lined. T Tauri star based on the IL, equivalent width from single epoch spectroscopy (2)..."," It is a G8-type star with a mass of $\sim$ $_{\odot}$, classified as a weak-lined T Tauri star based on the $_\alpha$ equivalent width from single epoch spectroscopy \citep{Alcala1993}."
 Subsequent photometric aud spectroscopic monitoring has indicated a strong variability of this line. which shows siguificaut changes iu its equivalent width. inteusitv. aud profile (2??)..," Subsequent photometric and spectroscopic monitoring has indicated a strong variability of this line, which shows significant changes in its equivalent width, intensity, and profile \citep[][]{GregorioHetem1992,
 Alcala1993, Schisano2009}."
" Tf the line is related to accretion episodes. then the average accretion rate is M = Lx ?AL,qe (1)."," If the line is related to accretion episodes, then the average accretion rate is $\dot{M}$ = $\times$ $^{-9}\,M_{\odot}/yr$ \citep{Schisano2009}."
 T Cha shows variable circumstellar extinetion with a nist frequent value of 421.7 ας according to ?.., T Cha shows variable circumstellar extinction with a most frequent value of $A_V$ mag according to \citet{Schisano2009}.
 The authors derive a disk extinction law characterized by Ry=5.5. which suggests the presence of large dust exains within the disk.," The authors derive a disk extinction law characterized by $R_V=5.5$, which suggests the presence of large dust grains within the disk."
 The age[m] of the source is variously estimated to be between LOMAIv according todifferent methods (72)..., The age of the source is variously estimated to be between Myr according to different methods \citep{Fernandez2008}.
 A colplete study of the € Cha association by ? provides au average age of My. while ? estimate a slightly older age (between OAIAIvr) based on the lithium couteut of he e Cha members.," A complete study of the $\epsilon$ Cha association by \citet{Torres2008} provides an average age of Myr, while \citet{daSilva2009} estimate a slightly older age (between Myr) based on the lithium content of the $\epsilon$ Cha members."
 Finally. a ανασα] evolution study of the (4 Cha custer. which probably belougs to the € Cha association. provides au age of NMyr (2)...," Finally, a dynamical evolution study of the $\eta$ Cha cluster, which probably belongs to the $\epsilon$ Cha association, provides an age of Myr \citep{Ortega2009}."
 For he purpose ofthis paper. we adopt au age of MM.," For the purpose of this paper, we adopt an age of Myr."
 The distance to the source. based on the Iipparcos xuallax. is )pedi15 ype.," The distance to the source, based on the Hipparcos parallax, is $\pm$ pc."
 A more reliable value of ppe was obtaiue uxue properanuotion studies (?7)..," A more reliable value of pc was obtained using proper-motion studies \citep[][]{Frink1998,Terranegra1999}."
 ? provided a kinematical distance of ppc for T Cha. aud an average value of 1084-9 ppc for the whole association.," \citet{Torres2008} provided a kinematical distance of pc for T Cha, and an average value of $\pm$ pc for the whole association."
 We adopted the latter value for this paper., We adopted the latter value for this paper.
 Finally. previous works based ou radial velocity (RV) and direct inagiug aud coronographlic techniques have not reported the presence of αν (stellar or very low-1ass) companion around T Cha (???)..," Finally, previous works based on radial velocity (RV) and direct imaging and coronographic techniques have not reported the presence of any (stellar or very low-mass) companion around T Cha \citep[][]{Schisano2009,Chauvin2010,Vicente2011}. ."
 The SAM observations allow us to fill the gap between between RW aud direct Huaging observations., The SAM observations allow us to fill the gap between between RV and direct imaging observations.
We iraced. through a hieh-resolution numerical experiment. the evolution of a core forming in a turbulent molecular cloud that is ina Jeans unstable and magnetically supercritical state.,"We traced, through a high-resolution numerical experiment, the evolution of a core forming in a turbulent molecular cloud that is in a Jeans unstable and magnetically supercritical state."
 The core is slowly contracting at (he central part with a subthermal speed., The core is slowly contracting at the central part with a subthermal speed.
 Due to ihe influence of accretion shocks and a nearby Gurbulent. flow. velocity fields in the outer region of the core can easily exceed (he thermal speed. and show not only the contracting but also sometimes expanding motions.," Due to the influence of accretion shocks and a nearby turbulent flow, velocity fields in the outer region of the core can easily exceed the thermal speed, and show not only the contracting but also sometimes expanding motions."
 Based on (he evolutionary profiles of densitv and eas and dust temperatures. we calculated the evolution of the abundance.," Based on the evolutionary profiles of density and gas and dust temperatures, we calculated the evolution of the $^{+}$ abundance."
 Once the central density of the core is larger than. 10em. the — abundance profiles peak at outer regions because of the depletion o£ at the central part of the core (see Figure 3).," Once the central density of the core is larger than $10^5~{\rm cm}^{-3}$ , the $^{+}$ abundance profiles peak at outer regions because of the depletion of $^{+}$ at the central part of the core (see Figure 3)."
 We also calculated the evolution of 3-2. 4—3. and CO 3—2 line profiles that are coupled to the density. kinetic temperature. and abundanuce structures.," We also calculated the evolution of $^{+}$ $-$ 2, $-$ 3, and $^{18}$ O $-$ 2 line profiles that are coupled to the density, kinetic temperature, and abundance structures."
 The most interesting result in (his work is that the — Hine profiles are dominantlv alfected by the velocity profiles at the outer regions., The most interesting result in this work is that the $^{+}$ line profiles are dominantly affected by the velocity profiles at the outer regions.
 This is the consequence of both the abundance ancl velocity profiles that have peaks at the outer regions., This is the consequence of both the $^{+}$ abundance and velocity profiles that have peaks at the outer regions.
 We showed that the molecular line profiles of a core. which results from the coupled velocily and chemical structures. vary wilh time.," We showed that the molecular line profiles of a core, which results from the coupled velocity and chemical structures, vary with time."
 To understand the formation process of cores. survey observations of many cores in different evolutionary states are crucial.," To understand the formation process of cores, survey observations of many cores in different evolutionary states are crucial."
 \lappine observations of individual cores with a high resolution enables us to know (he complex velocity structures of the cores., Mapping observations of individual cores with a high resolution enables us to know the complex velocity structures of the cores.
 The radial variation of the velocity field. as seen in BGS mav be an intrinsic part of star formation. aud the blue asvinmeiry is not always a direct signature of collapse in starless dense cores (Ixeto et al.," The radial variation of the velocity field as seen in B68 may be an intrinsic part of star formation, and the blue asymmetry is not always a direct signature of collapse in starless dense cores (Keto et al."
 2006)., 2006).
 We are verv grateful to Ted Berein for valuable comments., We are very grateful to Ted Bergin for valuable comments.
 This work was supported bv the IxXorea Science and Engineeringe Foundation under a cooperative agreementc» with the Astrophysical Research Canterfor the Structure and Evolution of the Cosmos., This work was supported by the Korea Science and Engineering Foundation under a cooperative agreement with the Astrophysical Research Canterfor the Structure and Evolution of the Cosmos.
In the Milky Way. stellar clusters form in dense regions located inside spiral arms.,"In the Milky Way, stellar clusters form in dense regions located inside spiral arms."
 When clusters survive. they remain connected with the parent arm for about 100 Myr (Dobbs Pringle 2010).," When clusters survive, they remain connected with the parent arm for about 100 Myr (Dobbs Pringle 2010)."
 Afterwards they decouple from it and are no longer useful as spiral structure tracers., Afterwards they decouple from it and are no longer useful as spiral structure tracers.
 Young star clusters have been used for half a century to probe the spiral structure of the Milky Way (MW)., Young star clusters have been used for half a century to probe the spiral structure of the Milky Way (MW).
 Historically. the first arms to be detected using young open clusters. were the Perseus arm in the second galactic quadrant. the Carina-Sagittarius in the fourth quadrant. and the Orion spur in which the Sun is embedded (Trumpler 1930).," Historically, the first arms to be detected using young open clusters, were the Perseus arm in the second galactic quadrant, the Carina-Sagittarius in the fourth quadrant, and the Orion spur in which the Sun is embedded (Trumpler 1930)."
 Nowadays. the picture we have of the MW spiral structure contains many more details (Russeil 2003. Efremov 201I. Lepine et al.," Nowadays, the picture we have of the MW spiral structure contains many more details (Russeil 2003, Efremov 2011, Lepine et al."
 2011). although a lively discussion is still ongoing as to how many major arms are present and if they are long-lived or transient (Grosbol et al.," 2011), although a lively discussion is still ongoing as to how many major arms are present and if they are long-lived or transient (Grosbol et al."
 2011)., 2011).
 In the last decade. young star clusters played a major role in improving our knowledge of the MW spiral structure. especially in the third Galactic quadrant.," In the last decade, young star clusters played a major role in improving our knowledge of the MW spiral structure, especially in the third Galactic quadrant."
 Moitinho et al. (, Moitinho et al. (
2005). Vazquez et al. (,"2005), Vazquez et al. ("
2008) and Carraro et al. (,2008) and Carraro et al. (
2010) identified for the first time in optical observations the outer Norma-Cygnus arm. and claritied the shape and interaction of the Orion and Perseus arms.,"2010) identified for the first time in optical observations the outer Norma-Cygnus arm, and clarified the shape and interaction of the Orion and Perseus arms."
 These studies made clear that young star clusters are powerful spiral tracers when it is possible to determine their distance and age with high contidence., These studies made clear that young star clusters are powerful spiral tracers when it is possible to determine their distance and age with high confidence.
 In particular. the authors stress how crucial deep (-band photometry is to pin down cluster reddening and hence obtain their," In particular, the authors stress how crucial deep $U$ -band photometry is to pin down cluster reddening and hence obtain their"
The standard cosmological model. (he ACDM. assumes the matter to be forever expanding in uünmost flat universe under the pressure of dark energv.,"The standard cosmological model, the$\Lambda$ CDM, assumes the matter to be forever expanding in utmost flat universe under the pressure of dark energy."
 In early Universe. before the recombination of electrons ancl protons ancl the formation of galaxies. the electron-proton plasma was assumed homogeneous.," In early Universe, before the recombination of electrons and protons and the formation of galaxies, the electron-proton plasma was assumed homogeneous."
 The formation of galaxies was considered as a consequence of gravitational instability of small initial density fluctuations 9p/pc10°. constrained by the results of WMAP mission (Hinshawetal.2009).," The formation of galaxies was considered as a consequence of gravitational instability of small initial density fluctuations $\delta \rho/\rho \sim 10^{-5}$, constrained by the results of WMAP mission \citep{wmap5}."
. In post-recombination epoch the gas of neutral atoms became utümnost (ransparent to photons., In post-recombination epoch the gas of neutral atoms became utmost transparent to photons.
 The CAIBR. with the present temperature 2.7. was no longer scattered by the matter (except for Che scattering on the free electrons of galaxies (5unvaev.andZeldovich 1970))). and therefore carries invaluable inlormation on (he early Universe.," The CMBR, with the present temperature 2.7K, was no longer scattered by the matter (except for the scattering on the free electrons of galaxies \citep{SZ1969s}) ), and therefore carries invaluable information on the early Universe."
 The ACDM model is an inflationary model. which depends on six parameters: the barion density O7. the cold darkmatter density Q7. a cosmological constant Q4. the spectral," The $\Lambda$ CDM model is an inflationary model, which depends on six parameters: the barion density $\Omega_b h^2$ , the cold darkmatter density $\Omega_c h^2$ , a cosmological constant $\Omega_\Lambda$, the spectral"
As ol Aor 2012. there are over TOO discovered exoplanets. and most of them are detected by he radial veocity (RV),"As of Apr 2012, there are over 700 discovered exoplanets, and most of them are detected by the radial velocity (RV)."
/exoplanets.org/.. RV precision of ] m-:s has been 'outiuely achieved (??) wih instrments such as HARPS (7) and HIRES (?).. which are ¢ross- echele spectrograj»hs.," RV precision of 1 $\rm{m}\cdot\rm{s}^{-1}$ has been routinely achieved \citep{Bouchy2009,Howard2010} with instruments such as HARPS \citep{Mayor2003} and HIRES \citep{Vogt1994}, which are cross-dispersed echelle spectrographs."
 WIille cross-clispersed echelle spectrograplis a‘e commonly used iu instruments for precision RV meast'ements. a method using a dispersed fixec delay interferometer (DEDI) has ofered an alteriative method (???)..," While cross-dispersed echelle spectrographs are commonly used in instruments for precision RV measurements, a method using a dispersed fixed delay interferometer (DFDI) has offered an alternative method \citep{Ge2006,Fleming2010,Lee2011}."
 Iu this method. a Michelso-type interferometer js used in combination with a imoclerate resolution spectrograph. RV signas are then extracted rour phase sliit of iuterference [riges of stellar absorption lines (????).. Tje.," In this method, a Michelson-type interferometer is used in combination with a moderate resolution spectrograph, RV signals are then extracted from phase shift of interference fringes of stellar absorption lines \citep{Erskine2000,Ge2002b,Ge2002,Erskine2003}."
 cletails about the DFDI theory and applicatious are «iscussed in ? and ?.., The details about the DFDI theory and applications are discussed in \citet{vanEyken2010} and \citet{Wang2011}.
 Instrumeut. acloptiο the DEDI method ias demonstrated acvalltages such as low cost. compact size aud imulti-objec," Instrument adopting the DFDI method has demonstrated advantages such as low cost, compact size and multi-object capability \citep{Ge2002,Ge2006,Fleming2010,Lee2011,Wisniewski2012}. ."
Further simplification of this equation is possible using the contravariant components of s: raising indices gives Lrom this we obtain where the second line used 0=8/2 and the third line used. Eq. 0362).,Further simplification of this equation is possible using the contravariant components of ${\bmath s}$: raising indices gives From this we obtain where the second line used $\Omega={\cal E}'/{\cal L}'$ and the third line used Eq. \ref{eq:wupt}) ).
 However. we also see that: ‘ere the second line used Eq. (35)):," However, we also see that: [Here the second line used Eq. \ref{eq:C020}) );"
 the third. line used. Eq. (123): , the third line used Eq. \ref{eq:slemma}) );
"the fourth line used ο=€'/L"" and the quotient rule: anc the fifth line used that w/w!=Qe C.", the fourth line used $\Omega={\cal E}'/{\cal L}'$ and the quotient rule; and the fifth line used that $w^\phi/w^t=\Omega={\cal E}'/{\cal L}'$ .]
 Equation (124)) leads to two major simplifications in Eq. (121))., Equation \ref{eq:biglemma}) ) leads to two major simplifications in Eq. \ref{eq:I2-intermed}) ).
 The term involving £ simplifies dramatically., The term involving $\xi^r$ simplifies dramatically.
 Also. using the first and third. lines of Ίσα. (124))," Also, using the first and third lines of Eq. \ref{eq:biglemma}) )"
 and €'2OU. we lind that Therefore. Eq. (121))," and ${\cal E}'=\Omega{\cal L}'$, we find that Therefore, Eq. \ref{eq:I2-intermed}) )"
" simplifies to A final level of simplification involves u,s."," simplifies to A final level of simplification involves $w^t{_{,r}}-s^t$."
 Using the explicit expressions. Eq. (15))," Using the explicit expressions, Eq. \ref{eq:wt}) )"
 for wá and Eq. (122)), for $w^t$ and Eq. \ref{eq:s-component}) )
 for s. we see that Finally. we consider 4;= Αμ.," for $s^t$, we see that This allows us to eliminate ${\bmath s}$ from our expression for $I_2$: Finally, we consider $I_3 = h^{r\alpha}w_\alpha$ ."
 This is most easilv computed. by explicit evaluation of the contravariant components using I5q. (105)):, This is most easily computed by explicit evaluation of the contravariant components using Eq. \ref{eq:habinv}) ):
 and This implies The terms involving € can be simplified using Eqs. (15)), and This implies The terms involving $\xi^r$ can be simplified using Eqs. \ref{eq:wt}) )
" and (16)). which simplifies them to we"")ove "," and \ref{eq:wphi}) ), which simplifies them to $\rmi m (\Omega_{\rm s}w^t-w^\phi)\xi^r$ ."
FurtherH using. «m=Qubos gives The other contributions to ὃ do not explicitly contain fi. so in order to prove gauge invariance we will need to eliminate f£ in favour of other variables.," Further using $w^\phi=\Omega w^t$ gives The other contributions to ${\cal S}$ do not explicitly contain $\mmu$, so in order to prove gauge invariance we will need to eliminate $\mmu$ in favour of other variables."
 Equation (28)) provides a convenient choice: it and the definitions of # and =tell us that We thus arrive at our final expression for Z5: We now substitute ἐν. fo. and £5 into Eq. (103)).," Equation \ref{eq:C002}) ) provides a convenient choice: it and the definitions of $\kappa$ and ${\cal Z}$tell us that We thus arrive at our final expression for $I_3$: We now substitute $I_1$, $I_2$, and $I_3$ into Eq. \ref{eq:FSM2}) ),"
" giving We may divide through by w on both sides. and cancel the terms involving 2v,,ο.wl)."," giving We may divide through by $w^t$ on both sides, and cancel the terms involving $2\nu_{,r}- \oomega_{,r}{\cal L}/({\cal E}-\oomega{\cal L})$."
 Collecting the remaining terms gives In general. this is nonzero.," Collecting the remaining terms gives In general, this is nonzero."
 However. there is one piece of information we have not used: that theresonant. amplitude is to be evaluated at the resonance location (I2)= 0. Le. When and only when we use this fact. we see that Eq. (136))," However, there is one piece of information we have not used: that theresonant amplitude is to be evaluated at the resonance location $D(R)=0$ , i.e. When – and only when – we use this fact, we see that Eq. \ref{eq:SM-simple}) )"
 vanishes., vanishes.
 That is.," That is,"
Condition I is related to the fact that we strive for a situation in which both the linear part as the nonlinear part of 1)) are expauded in the same basis functions.,Condition I is related to the fact that we strive for a situation in which both the linear part as the nonlinear part of \ref{eq: manakov}) ) are expanded in the same basis functions.
" Iu expermaüenuts. this situation can be obtained by optimizing the ""launchiug conditions”."," In experiments, this situation can be obtained by optimizing the “launching conditions”."
 Condition II is iutroduced by realizing that iu fiber optics oulv cuvelope solutions cau be measured., Condition II is introduced by realizing that in fiber optics only envelope solutions can be measured.
 Iu this paper. we search for modes 44 and «9 propagating with unequal velocity.," In this paper, we search for modes $u_{1}$ and $u_{2}$ propagating with unequal velocity."
 We therefore substitute the following solutions iuto l1» Iu Eq.(9)) it is used that :4=tdr and v2=p., We therefore substitute the following solutions into \ref{eq: manakov}) ): In \ref{soliton}) ) it is used that $z_{1}=t - \delta x$ and $z_{2}=t + \delta x$.
 For the solutions (9)) the left-hand side of Eq.(1)) is equal to: We can conclude that the linear part of 1)) is not mocdifving the structure othe exponential functions (9))., For the solutions \ref{soliton}) ) the left-hand side of \ref{eq: manakov}) ) is equal to: We can conclude that the linear part of \ref{eq: manakov}) ) is not modifying the structure of the exponential functions \ref{soliton}) ).
 This appears also to be the case for the nonlinear part., This appears also to be the case for the nonlinear part.
 If we substitute Eq.(9)) iuto the rnghbt-haud side of Eq.(1)) we obtain: It can be concluded by comparing Eq.(10)) aud. Eq.(113) that solutions of the type (9)) eusure that both the luear part and the nonlinear part of 1) can be expauded in the same basis functions., If we substitute \ref{soliton}) ) into the right-hand side of \ref{eq: manakov}) ) we obtain: It can be concluded by comparing \ref{leftt}) ) and \ref{rightt}) ) that solutions of the type \ref{soliton}) ) ensure that both the linear part and the nonlinear part of \ref{eq: manakov}) ) can be expanded in the same basis functions.
" Caven non-zero coefficients Ayy and Byy (A.BCIR ). all the other coefficients i4»,|4, aud D»,(1,5 of the solution (9)) are determined by the"," Given non-zero coefficients $\hat{A}_{1,0}$ and $\hat{B}_{1,0}$ $A,B \in I\!\!R$ ), all the other coefficients $\hat{A}_{2n+1,m}$ and $\hat{B}_{2n+1,m}$ of the solution \ref{soliton}) ) are determined by the"
colors.,colors.
 None of these objects are observed. (hus. we conclude all of the stars brighter than ihe AGB limit are true red supergiants.," None of these objects are observed, thus, we conclude all of the stars brighter than the AGB limit are true red supergiants."
 The right panel also shows the 3.6 TTIRGD., The right panel also shows the 3.6 TRGB.
 The value of the TRGB we adopt (M42 —6.22:0.2) was determined by inspecting the 3.6 luminosity function for sources detected at both 3.6 and 4.5jam., The value of the TRGB we adopt $_{3.6}$ $-$ $\pm$ 0.2) was determined by inspecting the 3.6 luminosity function for sources detected at both 3.6 and 4.5.
 The 3.6 luminosity function (the number of stars in each 0.2 magnitude bin at 3.6 jm) is shown in Figure 5.., The 3.6 luminosity function (the number of stars in each 0.2 magnitude bin at 3.6 ) is shown in Figure \ref{lum_fn}.
 We adopt this value of the TRGB based on the abrupt drop in detections at (hat magnitude., We adopt this value of the TRGB based on the abrupt drop in detections at that magnitude.
 Also. optically classified sub-TRGB red giants with blue —|4.5] colors are observed wilh Als magnitudes up to this value. but above this value only optically classified sub-TRGB red giants with very red colors are observed.," Also, optically classified sub-TRGB red giants with blue $-$ [4.5] colors are observed with $_{3.6}$ magnitudes up to this value, but above this value only optically classified sub-TRGB red giants with very red colors are observed."
" The IR. [Inxes of these objects are consistent with mass-losing AGBs (see 8??)) rather than sub-TRGBred giants. supporting our adopted value of Ma,47—6.2 for the 3.6 TRGD."," The IR fluxes of these objects are consistent with mass-losing AGBs (see \ref{AGB}) ) rather than sub-TRGBred giants, supporting our adopted value of $_{3.6}$ $-$ 6.2 for the 3.6 TRGB."
 This value is 0.2 magnitudes fainter than the value found lor the Large \lagellanic Cloud (LMC) (Alp -6.4: and 0.4 magnitudes lower (han (hat adopted for WLM (Jacksonetal.2007)., This value is 0.2 magnitudes fainter than the value found for the Large Magellanic Cloud (LMC) \citep[M$_{L^\prime}$ $-$ and 0.4 magnitudes lower than that adopted for WLM \citep{jac07}.
. It is unlikely that the difference in values of the TRGB between IC 1613 ancl WLM are actually so disparate. given their similar metallicities.," It is unlikely that the difference in values of the TRGB between IC 1613 and WLM are actually so disparate, given their similar metallicities."
 The difference in (hese values likely reflects the uncertainty in our adopted TRGB values., The difference in these values likely reflects the uncertainty in our adopted TRGB values.
 This uncertainty is not a major concern. however. because shifting the value of the TRGB by up to 0.5 magnitudes affects the detection statisties only slightly (see 82?)).," This uncertainty is not a major concern, however, because shifting the value of the TRGB by up to 0.5 magnitudes affects the detection statistics only slightly (see \ref{complete}) )."
 In general. the Mà; versus |3.6]— 4.5] CMD of IC! 1613 is very similar to that of WLM. with the only major dillerence being that there are significantly more sub-TRGB red giants detected in IC 1613 than in WLM because WLM is more distant.," In general, the $_{3.6}$ versus $-$ [4.5] CMD of IC 1613 is very similar to that of WLM, with the only major difference being that there are significantly more sub-TRGB red giants detected in IC 1613 than in WLM because WLM is more distant."
 As is the case of WLM. a small population of stars redward of (he main stellar distribution is detected. though fewer are seen in IC 1613 than in WLM.," As is the case of WLM, a small population of stars redward of the main stellar distribution is detected, though fewer are seen in IC 1613 than in WLM."
 As we diseuss in §??.. the infrared. fInxes ancl colors of these objects are consistent with mass-losing ACB stars.," As we discuss in \ref{mass_loss}, , the infrared fluxes and colors of these objects are consistent with mass-losing AGB stars."
 Figure 6 is the Mi versus [3.6]—[8.0] CMD for IC 1613., Figure \ref{ch4cmd} is the $_{8.0}$ versus $-$ [8.0] CMD for IC 1613.
 This CMD shows a narrow vertical feature with [8.0] -0 and another broad distribution of red objects with —12 < My « —1.5 ] < [3.6|-[8.0] < 4., This CMD shows a narrow vertical feature with $-$ $\sim$ 0 and another broad distribution of red objects with $-$ 12 $<$ $_{8.0}$ $<$ $-$ 7.5 1 $<$ $-$ [8.0] $<$ 4.
 These red objects are also the most luminous. reddest objects in the Ma; versus [3.6] [4.5] CMD. with typical colors between 0.5 and. 1.0.," These red objects are also the most luminous, reddest objects in the $_{3.6}$ versus $-$ [4.5] CMD, with typical colors between 0.5 and 1.0."
 Figure 7 shows the spatial distributions of different stellar tvpes based on their optical and IR. fluxes., Figure \ref{xy} shows the spatial distributions of different stellar types based on their optical and IR fluxes.
 The distributions of both red giants ancl AGB stars are verv smooth with no obvious concentrations aside from the radial stellar gradient., The distributions of both red giants and AGB stars are very smooth with no obvious concentrations aside from the radial stellar gradient.
 A thin gap is observed in the optical data. which is certainly an instrumental effect. though this is not mentioned in Udalskietal. (2001)..," A thin gap is observed in the optical data, which is certainly an instrumental effect, though this is not mentioned in \citet{uda01}. ."
 There are two conspicuous features in the blue objects. which outline the voung stellar distribution: a bar running across the midplane of thegalaxy. andl a large," There are two conspicuous features in the blue objects, which outline the young stellar distribution; a bar running across the midplane of thegalaxy and a large"
cm-?.,.
" While initially some heat is transported inwards, Brehmstrahlung cooling scales like n2; thus, with the addition of cooling, this density increase may enhance the cooling flow problem."," While initially some heat is transported inwards, Brehmstrahlung cooling scales like $n_e^2$; thus, with the addition of cooling, this density increase may enhance the cooling flow problem."
" Simultaneously, the MTT is driving a magnetic dynamo which acts to amplify the magnetic field."," Simultaneously, the MTI is driving a magnetic dynamo which acts to amplify the magnetic field."
 This is quantified as where the angle brackets denote a volume average over the sphere with r<Tmax., This is quantified as where the angle brackets denote a volume average over the sphere with $r<r_{\textrm{max}}$.
" For Run A1, the magnetic energy density is amplified by a factor of 18.4 from the initial energy density."," For Run A1, the magnetic energy density is amplified by a factor of 18.4 from the initial energy density."
" There is actually somewhat more magnetic dynamo action, as some of the amplified flux is transported by penetrative convection to r>Tmax and is not included in this average amplification factor."," There is actually somewhat more magnetic dynamo action, as some of the amplified flux is transported by penetrative convection to $r>r_{\textrm{max}}$ and is not included in this average amplification factor."
 The saturation properties of this run and all other simulations are given in Table 4.., The saturation properties of this run and all other simulations are given in Table \ref{tab:satprop}.
" This magnetic field increase is significant, but clearly not entirely sufficient to raise the primordial field to its present-day value."," This magnetic field increase is significant, but clearly not entirely sufficient to raise the primordial field to its present-day value."
" Certainly, much of the increase comes from flux-freezing during gravitational collapse, which is not included here."," Certainly, much of the increase comes from flux-freezing during gravitational collapse, which is not included here."
" In addition, smaller-scale field increases could come from purely kinetic processes such as the Alfvénn wave cascadeor other kinetic process ??).. "," In addition, smaller-scale field increases could come from purely kinetic processes such as the Alfvénn wave cascadeor other kinetic process \citep[e.g.,][]{sc06a, sc06b}."
"The large scale turbulence can serve as energy(e.g., input for this process.", The large scale turbulence can serve as energy input for this process.
 We proceed to discuss the more realistic fiducial case of run A2., We proceed to discuss the more realistic fiducial case of run A2.
" Here, instead of a fixed conductivity as in Al, we utilize the full variable Spitzer conductivity as given in Eqn. [26]]."," Here, instead of a fixed conductivity as in A1, we utilize the full variable Spitzer conductivity as given in Eqn. \ref{eqn:clust:conductivity}] ]."
 We begin by examining the thermal evolution of A2 in Figure 3.., We begin by examining the thermal evolution of A2 in Figure \ref{fig:clust:A2-temp}. .
The first observational evidence of PF stellar nucleosvnthesis was reported by Jorissen. snuth Lambert (1992. herealter JSL).,"The first observational evidence of $^{19}$ F stellar nucleosynthesis was reported by Jorissen, Smith Lambert (1992, hereafter JSL)."
 These authors derived F enhancements up to a factor 50 solar in a sample of Galactic AGB stars. and found a correlation between {his enhancement and the C/O ratio.," These authors derived F enhancements up to a factor 50 solar in a sample of Galactic AGB stars, and found a correlation between this enhancement and the C/O ratio."
 Since the C/O is expected (ο increase as a consequence ol third dredge up (EDU) episodes during the AGB phase (e.g. Busso et al., Since the C/O is expected to increase as a consequence of third dredge up (TDU) episodes during the AGB phase (e.g. Busso et al.
 1999). this," 1999), this"
ab the long wavelength end.,at the long wavelength end.
 The presence of strong emission lines coupled with the low flux level and the steeply rising continuum are not characteristic of the FUV spectra of nova-like variables in their high optical brightness states unless (he inclination angle is high., The presence of strong emission lines coupled with the low flux level and the steeply rising continuum are not characteristic of the FUV spectra of nova-like variables in their high optical brightness states unless the inclination angle is high.
 We note that the emission lines aud continuum slope in BIX Lvn's specüirum are strikinely similar in appearence to the IST and HUT spectra of the dwarf nova SS Cyeni in quiescence (see Fig.l in Long et al.2005) as well as the FUV spectrum of V794 Αα in its high state (Godon et aL2007)., We note that the emission lines and continuum slope in BK Lyn's spectrum are strikingly similar in appearence to the HST and HUT spectra of the dwarf nova SS Cygni in quiescence (see Fig.4 in Long et al.2005) as well as the FUV spectrum of V794 Aql in its high state (Godon et al.2007).
 The F28x50LP magnitude from the acquisition exposure was 15.3. roughly consistent with a high state but slightly fainter than the normal visual magnitude range of 14.6 to 14.7.," The $\times$ 50LP magnitude from the acquisition exposure was 15.3, roughly consistent with a high state but slightly fainter than the normal visual magnitude range of 14.6 to 14.7."
 The STIS spectrum of V151 νο reveals moderately strong absorption features seen against a continuum rising toward shorter wavelengths., The STIS spectrum of V751 Cygni reveals moderately strong absorption features seen against a continuum rising toward shorter wavelengths.
 The FUV continuum even after reddening. has a relatively flat slope while the Lya profile is rather narrow.," The FUV continuum even after de-reddening, has a relatively flat slope while the $\alpha$ profile is rather narrow."
 This poses a considerable challenge in finding an accretion disk spectral energy distribution which satisfies both the continuum slope and vields the observed Lvo line width in a single disk moclel., This poses a considerable challenge in finding an accretion disk spectral energy distribution which satisfies both the continuum slope and yields the observed $\alpha$ line width in a single disk model.
 The narrow width of Lya should be associated with a much steeper continuum slope than is observed in the spectra., The narrow width of $\alpha$ should be associated with a much steeper continuum slope than is observed in the spectra.
 Our LST STIS spectrum confirms the identification of Ile II (1640) in absorption although the equivalent width of the feature is siialler (han seen in the IVE spectrum (see Section 1.1 above)., Our HST STIS spectrum confirms the identification of He II (1640) in absorption although the equivalent width of the feature is smaller than seen in the IUE spectrum (see Section 1.1 above).
 Thed C IV (1550) absorption lines reveals a hint of P Cvgni structure probably indicating wind noutflow at the time the STIS spectrum was obtained., Thed C IV (1550) absorption lines reveals a hint of P Cygni structure probably indicating wind noutflow at the time the STIS spectrum was obtained.
 The STIS spectrum of V380 Oph is dominated by strong broad emission leatures superimposed on a continuum rising toward shorter wavelengths For V380 Oph., The STIS spectrum of V380 Oph is dominated by strong broad emission features superimposed on a continuum rising toward shorter wavelengths For V380 Oph.
 We found the same inconsistency between. (he observed. width of the Ίωνα profile and the observed continuuun slope., We found the same inconsistency between the observed width of the $\alpha$ profile and the observed continuuum slope.
 Several weak. sharp absorption features are most likely of interstellar origin.," Several weak, sharp absorption features are most likely of interstellar origin."
 For V380 Oph. we found (he same inconsistency between the observed. width of the Lya prolile aud the observed continuuunm slope.," For V380 Oph, we found the same inconsistency between the observed width of the $\alpha$ profile and the observed continuuum slope."
 This led us to first explore the effect. of interstellar reddening on the spectrum., This led us to first explore the effect of interstellar reddening on the spectrum.
 In table 3. we have listed the strongest spectral features detected for each svstem.," In table 3, we have listed the strongest spectral features detected for each system."
otted iu Fie. 1..,"plotted in Fig. \ref{fig:vgvsk},"
 which shows that the behavior of the wo curves is different at small aud iuterinediate inuubers: m the horizoutal direction the group velocity relmaius snall up to &~0.7 and then increases rapidly to je asvinptotic value of unity. aud in the vertical direction ie erowth beeins steeply at &=0 and then coutinucμα nore eracdually to the asviuptotic vali," which shows that the behavior of the two curves is different at small and intermediate numbers: in the horizontal direction the group velocity remains small up to $ k \sim 0.7 $ and then increases rapidly to the asymptotic value of unity, and in the vertical direction the growth begins steeply at $k= 0$ and then continues more gradually to the asymptotic value."
 The evolution of the distribution of wavenuubers in PAoe Is described by Eqs. (139), The evolution of the distribution of wavenumbers in space is described by Eqs. \ref{eqn1}) )
" and (11). with the eroup velocity v, eiven by Eqs. (15))"," and \ref{eqn2}) ), with the group velocity $\vg$ given by Eqs. \ref{vgz}) )"
 and (16))., and \ref{vgp}) ).
 Thex equations can be solved for the compoueuts A. and Ay oei the vertical aud horizoutal directions as functions of Djfonud w/t: the soluious are shown in Fig. 2.., These equations can be solved for the components $k_z$ and $k_\perp$ in the vertical and horizontal directions as functions of $z/t$ and $x/t$; the solutions are shown in Fig. \ref{fig:vg}.
 The two panels reflect the different. behavior. described above. of the eroup velocity as a function of wavenumber: in both cases. sluall wavenmubers. which have low group velocity. are found close to the origin. while large wavemmubers. which have group velocity close to the speed of sound. are found near the pulse.," The two panels reflect the different behavior, described above, of the group velocity as a function of wavenumber; in both cases, small wavenumbers, which have low group velocity, are found close to the origin, while large wavenumbers, which have group velocity close to the speed of sound, are found near the pulse."
 At intermediate distances. the ΠΠtion of À valies ds much broader in the vertical than iu the horizontal direction.," At intermediate distances, the distribution of $k$ values is much broader in the vertical than in the horizontal direction."
 This las uuediate consequences for the nature of fjio. solution: since the value of the wavemuuber found at a particular position nuplies also the waveleleusth of the oscillation. oue may expect (see Fie.," This has immediate consequences for the nature of the solution: since the value of the wavenumber found at a particular position implies also the wavelelength of the oscillation, one may expect (see Fig."
 6 below) that in the horizoutal direction the oscillations behind the frout have almost coustaut wavelength. while in the vertical direction the wavelength eradually increases from the head of the wave towards tlic Origin.," \ref{fig:acvsxz} below) that in the horizontal direction the oscillations behind the front have almost constant wavelength, while in the vertical direction the wavelength gradually increases from the head of the wave towards the origin."
" Combining the plase velocity from the dispersion relation with the cdistribition of wavenunbers obtained from the evolution equajous (13)) aud (11)) one cau follow the propagation o© points of constant phase. for exanirple of a ανα,"," Combining the phase velocity from the dispersion relation with the distribution of wavenumbers obtained from the evolution equations \ref{eqn1}) ) and \ref{eqn2}) ) one can follow the propagation of points of constant phase, for example of a maximum."
 When a particular maxima is located close to the origi ων= 0) it is characterized by low values of the wavenunber. aud thus high phase velocity.," When a particular maximum is located close to the origin $x,z=0$ ) it is characterized by low values of the wavenumber, and thus high phase velocity."
 As it μονος away from the origin. its local wavenunuber as well as its frequency increase (in the vertical- direction.HB 47D=A7|E 1). whileB itsB propagation- velocity decreases (Fig. 2).," As it moves away from the origin, its local wavenumber as well as its frequency increase (in the vertical direction, $\omega^2=k_z^2+1$ ), while its propagation velocity decreases (Fig. \ref{eq:vel}) )."
 Thus. as oa particular Παππια travels away from the origin and towards the read of the wave it is ducreasinely characterized by high-requency components (which become important for shock ornation iu the uoulinear τοσο).," Thus, as a particular maximum travels away from the origin and towards the head of the wave it is increasingly characterized by high-frequency components (which become important for shock formation in the nonlinear regime)."
 Fie., Fig.
 preseuts this vchavior i more detail it shows the phase velocity of wee consecutive nixiua as fuuctious of distance from ie source for the vertical (upper panel) aud rorizoutal ower pane) diectious., presents this behavior in more detail; it shows the phase velocity of three consecutive maxima as functions of distance from the source for the vertical (upper panel) and horizontal (lower panel) directions.
" In the horizoutal direction the shase velocity. is practically equal to he soiud speed jiearlv throughout the whole region. aud for all παλια,"," In the horizontal direction the phase velocity is practically equal to the sound speed nearly throughout the whole region, and for all maxima."
 Iu the vertical direction the behavior is wore complicated: Except for the head of the wave (the first maxiumuu in Fig. 3..," In the vertical direction the behavior is more complicated: Except for the head of the wave (the first maximum in Fig. \ref{fig:vpvsxz},"
 upper panel). which propagates at the sound speed. the later maxima have high phase velocity over oewcreasinely extended παςit ranges.," upper panel), which propagates at the sound speed, the later maxima have high phase velocity over increasingly extended height ranges."
 They therefore travel at increased phase velocity* (Fig. 6..," They therefore travel at increased phase velocity (Fig. \ref{fig:acvsxz},"
" upper panel: see also IWRBAL 22). especial~ near the origin (where xoas lo> 0) until they apxoach the head of the wave,"," upper panel; see also KRBM 2), especially near the origin (where $v_{\rm ph}\rightarrow\infty$ as $z\rightarrow0$ ), until they approach the head of the wave."
 As a consequence of this littevence in behavior between the two directions. the surfaces of constant plase behind the," As a consequence of this difference in behavior between the two directions, the surfaces of constant phase behind the"
dwarts.,dwarfs.
 In addition. loug timescale. high signal-to-noise observations can be used to check for possible periodicity in the radio emission. as has been uncovered for the L3.5 dwart 2500036|18 (Bergeretal.2005).," In addition, long timescale, high signal-to-noise observations can be used to check for possible periodicity in the radio emission, as has been uncovered for the L3.5 dwarf 0036+18 \citep{brr+05}."
. If the periodicity is in fact related to a close-in companion. Which excites the dvuameo by tidal or magnetic interactions. this iav be a ubiquitous feature of the active objects.," If the periodicity is in fact related to a close-in companion, which excites the dynamo by tidal or magnetic interactions, this may be a ubiquitous feature of the active objects."
 Finally. these same objects should be observed siunultaueouslv in the radio. N-ravs. aud Πα in order to directly measure the correlation. or lack thereof. betweeu these activity indicators.," Finally, these same objects should be observed simultaneously in the radio, X-rays, and $\alpha$ in order to directly measure the correlation, or lack thereof, between these activity indicators."
 This is necessary in order to trace the origin of the shift iu the radio/X-rav correlation. and in order to trace the evolution of flares as the release of maeuctic stresses. evident iu the radio. heats up the corona (N-ravs) and chromosphere (Ia).," This is necessary in order to trace the origin of the shift in the radio/X-ray correlation, and in order to trace the evolution of flares as the release of magnetic stresses, evident in the radio, heats up the corona (X-rays) and chromosphere $\alpha$ )."
 While the απο! observations allow us to address in a statistical manner the overall treuds observed with cach teclinique. oulv simultancous observations can provide iusieht iuto the production aud evolution of flares.," While the current observations allow us to address in a statistical manner the overall trends observed with each technique, only simultaneous observations can provide insight into the production and evolution of flares."
 As the λος catastrophic eveuts in the atinosplieres of dwarf stars. stich events will undoubtedly shed light on the structure of the fields. their streneths. and the details of the eunergv dissipation process. all of which will provide observatioua coustraits on the dvnamo mechanisia in dwarf stars.," As the most catastrophic events in the atmospheres of dwarf stars, such events will undoubtedly shed light on the structure of the fields, their strengths, and the details of the energy dissipation process, all of which will provide observational constraints on the dynamo mechanism in dwarf stars."
 Research has benefitted from the M. L. aud T. dwarf compcudit housed at DwarfArchives.org and maiutainec by Chris Coliuo. Davy Kirkpatrick. and Adam Bureasser.," Research has benefitted from the M, L, and T dwarf compendium housed at DwarfArchives.org and maintained by Chris Gelino, Davy Kirkpatrick, and Adam Burgasser."
 This work has ade use of the SIMDBAD database. operated at CDS. Strasbourg. France.," This work has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
 E.D. is supporte is supported by NASA through IIubble Fellowship eraut IIST-01171.01 awarded bv the Space Telescope Scieuce Iustitute. which is operated by AURA. Inc.. for NASA under contract NAS 5-26555.," E.B. is supported is supported by NASA through Hubble Fellowship grant HST-01171.01 awarded by the Space Telescope Science Institute, which is operated by AURA, Inc., for NASA under contract NAS 5-26555."
that for suitably chosen configurations a deconfinemenut transition iu the interior can occur upon spin-down.,that for suitably chosen configurations a deconfinement transition in the interior can occur upon spin-down.
 Tn order to demonstrate the cousistency of our perturbative approach. we show in Fig.," In order to demonstrate the consistency of our perturbative approach, we show in Fig."
 1 the values of the expansion parameter for the maximally attainable rotation frequencies of stationary rotating objects (Qj ος) as a function of the ceutral density characterizing the configuration., \ref{fig4} the values of the expansion parameter for the maximally attainable rotation frequencies of stationary rotating objects $\Omega_K/\Omega_0$ ) as a function of the central density characterizing the configuration.
 Tn Fig., In Fig.
 5 we show the critical regions of tle plase transition in the inner structure of the star configuration as well as the equatorial aud polar radii iu the plane of aneular velocity © versus distance from the ceuter of the star., \ref{fig5} we show the critical regions of the phase transition in the inner structure of the star configuration as well as the equatorial and polar radii in the plane of angular velocity $\Omega$ versus distance from the center of the star.
 It is obvious that with the increase of the aneular velocity the star is deforming its shape., It is obvious that with the increase of the angular velocity the star is deforming its shape.
 The maximal cecentricitics of the configurations with Np=L3N.. Np=LosXN. aud Np=LaΑν are e(Oκ)=0.7603. ες)=0.7655 aud e(Oj)=0.7659. respectively.," The maximal eccentricities of the configurations with $N_B=1.3~N_\odot$ , $N_B=1.55~N_\odot$ and $N_B=1.8~N_\odot$ are $\epsilon(\Omega_K)=0.7603$, $\epsilon(\Omega_K)=0.7655$ and $\epsilon(\Omega_K)=0.7659$, respectively."
 Duc to the changes of the ceutra density the quark core could disappear above a critical augular velocity., Due to the changes of the central density the quark core could disappear above a critical angular velocity.
 Iu Fig., In Fig.
 6 we display the dependence of the moment of inertia ou the angular velocity for configurations with the same total barvon munber Vp=1.55N. together with the different contributions to the total change of the moment of inertia., \ref{fig6} we display the dependence of the moment of inertia on the angular velocity for configurations with the same total baryon number $N_B=1.55~N_\odot$ together with the different contributions to the total change of the moment of inertia.
 As it is shown the most iuportaut contributions come frou the mass redistribution auk the shape deformation., As it is shown the most important contributions come from the mass redistribution and the shape deformation.
" The relativistic contributions due to field ancl rotational euerevOo, are less naportanut.", The relativistic contributions due to field and rotational energy are less important.
 Iu the same Fig., In the same Fig.
 6 we show the decrease of the spherical moment of inertia due to the decrease of the ceutral density for high aneular velocities which tends to partially compensate the further increase of the total moment of inertia for arec ., \ref{fig6} we show the decrease of the spherical moment of inertia due to the decrease of the central density for high angular velocities which tends to partially compensate the further increase of the total moment of inertia for large $\Omega$.
 There is no dramatic chauge in the slope of £(Q) at ent=2.77 Iz., There is no dramatic change in the slope of $I(\Omega)$ at $\Omega_{\rm crit}=2.77$ kHz.
 Fie., Fig.
 7 shows the dependence of the moment of inertia asa function of the augular velocity., \ref{fig7} shows the dependence of the moment of inertia as a function of the angular velocity.
 It is demonstrated that the behavior of £(Q) for a given total umber of barvous Vp strouely depeuds on the presence of a pure quark matter core in the ceuter of the star., It is demonstrated that the behavior of $I(\Omega)$ for a given total number of baryons $N_B$ strongly depends on the presence of a pure quark matter core in the center of the star.
 If the core does already. exist or it does not appear when the augularo velocity increascs up to the παπα value Όμως then the second. order derivative of the moment of inertia I(9) does not change itssign., If the core does already exist or it does not appear when the angular velocity increases up to the maximum value $\Omega_{\rm max}$ then the second order derivative of the moment of inertia $I(\Omega)$ does not change itssign.
" For he configuration with Np=LOIN, the critical value for the occurrence of the", For the configuration with $N_B=1.55~N_\odot$ the critical value for the occurrence of the
 During the present decade the number of confirmed. solar-like pulsators — those with acoustic modes excited by turbulent motions in the near-surface convection (e.g.?.andreferencetherein) — has increased enormously thanks. first. to the growing number of ground-based observing campaigns (e.g.??). and second. to the high-precision photometry measurements provided by space instrumentation such as WIRE (Wide-FieldInfraredExplorer.e.g.?).. MOST (MicrovariabilityandOscillationsofSTars.2). and CoRoT (?)..," During the present decade the number of confirmed solar-like pulsators – those with acoustic modes excited by turbulent motions in the near-surface convection \citep[e.g.][and reference therein]{2004SoPh..220..137C} – has increased enormously thanks, first, to the growing number of ground-based observing campaigns \citep[e.g.][]{2007CoAst.150..106B,2008ApJ...687.1180A}, and second, to the high-precision photometry measurements provided by space instrumentation such as WIRE \citep[Wide-Field Infrared Explorer, e.g.][]{2007A&A...461..619B}, MOST \citep[Microvariability and Oscillations of STars,][]{2003PASP..115.1023W} and CoRoT \citep{2006ESASP.624E..34B}."
 The latter has been providing data with an unprecedented quality both in terms of photometric precision and in terms of uninterrupted observation lengths., The latter has been providing data with an unprecedented quality both in terms of photometric precision and in terms of uninterrupted observation lengths.
 CoRoT has already observed several main-sequence solar-like pulsators (?) while ithas enabled to resolve the individual modes of the oscillations spectra of several F stars (2???) and G stars (?.Bal-lotetal.in prep.)..," CoRoT has already observed several main-sequence solar-like pulsators \citep{2008Sci...322..558M} while it has enabled to resolve the individual modes of the oscillations spectra of several F stars \citep{2008A&A...488..705A,2009A&A...506...51B,2009A&A...506...41G,2009A&A...506...33M} and G stars \citep[][Ballot et al. in prep.]{2010arXiv1003.4368D}."
 At least. it allowed to derive a large spacing for the faintest targets (??)..," At least, it allowed to derive a large spacing for the faintest targets \citep{2009A&A...506...41G,2009A&A...506...33M}."
 The measurement of these seismic parameters already offer a valuable tool for accurate determinations of radii (e.g.?) and ages (e.g.?) of stars. which are specially interesting for better understand stellar evolution as well as to characterise stars hosting planets.," The measurement of these seismic parameters already offer a valuable tool for accurate determinations of radii \citep[e.g.][]{stello09} and ages \citep[e.g.][]{2008arXiv0810.2440C} of stars, which are specially interesting for better understand stellar evolution as well as to characterise stars hosting planets."
 All of these CoRoT observations are the starting point for a better understanding of the structure (??) and the surface dynamics (???) of this class of stars.," All of these CoRoT observations are the starting point for a better understanding of the structure \citep{2009A&A...506..175P,2009Ap&SS.tmp..241D} and the surface dynamics \citep{2009A&A...506..167L,2009arXiv0910.4027S,2009arXiv0910.4037S} of this class of stars."
 The Kepler mission. successfully launched in March. 7. 2009 (2).. will also contribute to this field by observing stars on very long runs (4 years).," The Kepler mission, successfully launched in March 7, 2009 \citep{2009IAUS..253..289B}, will also contribute to this field by observing stars on very long runs (4 years)."
 The quality. of the first. data on stars showing solar-like oscillations (2222) promises the asteroseismology a bright future on the study of stellar interiors and dynamical processes (22?)..," The quality of the first data on stars showing solar-like oscillations \citep{2010ApJ...713L.176B, 2010ApJ...713L.169C, 2010ApJ...713L.187H, 2010ApJ...713L.182S} promises the asteroseismology a bright future on the study of stellar interiors and dynamical processes \citep[][]{2009arXiv0911.4629C,2009arXiv0912.0817S,2009A&A...506..811M}."
 In this paper we present results about à star recently observed by CoRoT. HD 170987 (or HIP 90851).," In this paper we present results about a star recently observed by CoRoT, HD 170987 (or HIP 90851)."
 This target is a well-known double star where components are separated by 0.7” (e.g.2).., This target is a well-known double star where components are separated by $0.7\arcsec$ \citep[e.g.][]{2002yCat.1274....0D}.
 The main star is a F5 dwarf star with a magnitude my ranging from 7.4 to 7.7 in the literature. while the second component has a magnitude around 8.5.," The main star is a F5 dwarf star with a magnitude $m_V$ ranging from 7.4 to 7.7 in the literature, while the second component has a magnitude around 8.5."
 We start by reporting the latest spectroscopic results observed by the NARVAL spectrograph (in Sect., We start by reporting the latest spectroscopic results observed by the NARVAL spectrograph (in Sect.
 2). which shows that this star is very similar to Procyon (e.g.2).," 2), which shows that this star is very similar to Procyon \citep[e.g.][]{allende02}."
. Then we describe the observations done with CoRoT during 149 days and the interpolation done in the data gaps of the light curve (Sect., Then we describe the observations done with CoRoT during 149 days and the interpolation done in the data gaps of the light curve (Sect.
 3)., 3).
 In Sect., In Sect.
 4 we infer the surface-rotation period of the star from the detailed analysis of the low-frequency region of the power spectrum and then. we obtain the global properties of the acoustic modes and of the star. respectively in Sect.," 4 we infer the surface-rotation period of the star from the detailed analysis of the low-frequency region of the power spectrum and then, we obtain the global properties of the acoustic modes and of the star, respectively in Sect."
 5 and Sect., 5 and Sect.
 6., 6.
 We finish in Sect., We finish in Sect.
 7 with a discussion of the results and in Sect., 7 with a discussion of the results and in Sect.
 8 with the conclusions of the paper., 8 with the conclusions of the paper.
 We have observed HD 170987 using the NARVAL spectrograph on the 2-m class Bernard Lyot Telescope at the Pic du Midi Observatory., We have observed HD 170987 using the NARVAL spectrograph on the 2-m class Bernard Lyot Telescope at the Pic du Midi Observatory.
 We acquired one spectrum on each night of 2009 July 7. 8 10. 11 and 13.," We acquired one spectrum on each night of 2009 July 7, 8 10, 11 and 13."
 The spectra were co- to obtain a signal-to-noise ratio m the continuum of, The spectra were co-added to obtain a signal-to-noise ratio in the continuum of
with the near close packing of the galaxies al 2;.,with the near close packing of the galaxies at $z_i$.
 The conversion [rom model gealactocentric to heliocentric velocities depends on the circular velocity of the local standard of rest., The conversion from model galactocentric to heliocentric velocities depends on the circular velocity of the local standard of rest.
 The 47 minimization brings the circular velocity lo vy. =1., The $\chi^2$ minimization brings the circular velocity to v_c =.
 Table 2 compares measured and model proper motions of LMC. M33 and 1C10.," Table 2 compares measured and model proper motions of LMC, M33 and IC10."
 The largest discrepancy is the motion of ICIO in declination. at three (mes the measurement uncertainty.," The largest discrepancy is the motion of IC10 in declination, at three times the measurement uncertainty."
 This is arguably acceptable within the approximations of the model. and considering the diflienlty of the measurement.," This is arguably acceptable within the approximations of the model, and considering the difficulty of the measurement."
 Table + lists heliocentric velocities normal to the line of sight in the directions of increasing right ascension aud declination. and (he corresponding angular velocities. for all the LG galaxies.," Table 4 lists heliocentric velocities normal to the line of sight in the directions of increasing right ascension and declination, and the corresponding angular velocities, for all the LG galaxies."
 The mass-to-light ratios m/Ly in Table 4 are based on the 2\JASS Ix-band. absolute magnitudes Ly in the Local Universe Catalog., The mass-to-light ratios $m/L_K$ in Table 4 are based on the 2MASS K-band absolute magnitudes $L_{\rm K}$ in the Local Universe Catalog.
 These magnitudes are reliable for the larger ealaxies. quite uncertain lor the least. luminous ones.," These magnitudes are reliable for the larger galaxies, quite uncertain for the least luminous ones."
 The smaller galaxies were arbitrarily assigned nominal masses [or the computation of 4? derived [rom m/Ly=50., The smaller galaxies were arbitrarily assigned nominal masses for the computation of $\chi^2$ derived from $m/L_{\rm K}=50$.
 Where that mass is oo small to affect the solution the final mass nevertheless is changed by the insistence ol the computation on small but nonzero (rial mass shifts in the approach to a minimum of > V. but the change isB small.," Where that mass is too small to affect the solution the final mass nevertheless is changed by the insistence of the computation on small but nonzero trial mass shifts in the approach to a minimum of $\chi^2$ , but the change is small."
 :The ~19 final. values Οι50 signilv.ae masses (hat seem {ο be too small to matter., The $\sim 19$ final values $m/L_{\rm K}\sim 50$ signify masses that seem to be too small to matter.
 The possible significance of the curious values of some of the other asses is considered in Section ??.., The possible significance of the curious values of some of the other masses is considered in Section \ref{sec:discussion}.
 Figures 3 to 5 show orthogonal views of the model orbits plotted in comoving supergalactic coordinates., Figures \ref{Fig:3} to \ref{Fig:5} show orthogonal views of the model orbits plotted in comoving supergalactic coordinates.
 The center of mass of the LG galaxies plus the external actors is al rest and (he origin is at (he present position of MW., The center of mass of the LG galaxies plus the external actors is at rest and the origin is at the present position of MW.
 The labels near initial positions al expansion factor 1+2;=10 are keved to names in the first two columns in Tables 1. 3. and 4.," The labels near initial positions at expansion factor $1+z_i=10$ are keyed to names in the first two columns in Tables 1, 3, and 4."
 The orbits of the two dominant Local Group galaxies. MW. and M31. ave plotted in red. and the four external actors that are supposed to eive a phenomenological description of the ellect of mass outside 1.5 kpc are shown in blue.," The orbits of the two dominant Local Group galaxies, MW and M31, are plotted in red, and the four external actors that are supposed to give a phenomenological description of the effect of mass outside $1.5$ kpc are shown in blue."
 The orbits of the four mislit galaxies with the poorest fits to catalog redshilts and distances are plotted as the dashed black curves. ancl the orbits of the other LG galaxies areshown as the solid black curves.," The orbits of the four misfit galaxies with the poorest fits to catalog redshifts and distances are plotted as the dashed black curves, and the orbits of the other LG galaxies areshown as the solid black curves."
"however, it is Alice (Bob) who measures first and thus causes the collapse of the other particle.","however, it is Alice (Bob) who measures first and thus causes the collapse of the other particle."
 So the question is solved in a very curious way., So the question is solved in a very curious way.
 Different observers plainly on the moment at which the co-collapse occurs (see Figure but they all see an order of events which is in perfect B)concordance with causality., Different observers plainly on the moment at which the co-collapse occurs (see Figure \ref{fig:EPR_boost}) ) but they all see an order of events which is in perfect concordance with causality.
" So in the experiment at hand, the (hard) question on the wave function reduction of the entangled particle occurs somehow seems to be irrelevant."," So in the experiment at hand, the (hard) question on the wave function reduction of the entangled particle occurs somehow seems to be irrelevant."
" However, one might be suspicious."," However, one might be suspicious."
 Maybe things do not always work out so nicely., Maybe things do not always work out so nicely.
 This is precisely where the delayed choice experiments come in., This is precisely where the delayed choice experiments come in.
" We briefly review these experiments, and show that essentially the above question/problem re-appears, but now in a matured version."," We briefly review these experiments, and show that essentially the above question/problem re-appears, but now in a matured version."
" Originally, the notion of ‘delayed choice’ arose in a set of thought experiments, devised by Wheeler[6]."," Originally, the notion of `delayed choice' arose in a set of thought experiments, devised by Wheeler."
". Since then, several variants have found experimental realization[L4]."," Since then, several variants have found experimental realization."
". For concreteness, we shall restrict our attention to one specific (slightly alternative but more instructive) realization, the eraser."," For concreteness, we shall restrict our attention to one specific (slightly alternative but more instructive) realization, the."
". A discussion of another (more literal, but less instructive) variant can be found in Appendix A. 'The delayed choice quantum eraser experiment was first realized in[5]."," A discussion of another (more literal, but less instructive) variant can be found in Appendix A. The delayed choice quantum eraser experiment was first realized in."
 The details are shown inFigure Bl., The details are shown inFigure \ref{fig:eraser}.
" The outcome is as follows: signal photons for which the corresponding idler photon later reveals which-path information, do not show an interference pattern."," The outcome is as follows: signal photons for which the corresponding idler photon later reveals which-path information, do not show an interference pattern."
" Their detection rates are precisely those of collapsed, single slit paths."," Their detection rates are precisely those of collapsed, single slit paths."
" Signal photons for which the idler does not reveal any path-information, form an untouched interference pattern."," Signal photons for which the idler does not reveal any path-information, form an untouched interference pattern."
 So interference at Do only occurs for events where the idler photon is detected at D4 or D»., So interference at $D_0$ only occurs for events where the idler photon is detected at $D_1$ or $D_2$ .
use of the sly constraint discussed in Section 4. ie.. 9=0.,"use of the $A_V$ constraint discussed in Section 4, i.e., $\beta=0$."
 The RAIS cliflerence between the estimated ancl spectroscopic values of logZur for that case was 0.18. corresponding to a percentage error of in the estimated value of Thy.," The RMS difference between the estimated and spectroscopic values of $\log T_{\rm eff}$ for that case was 0.18, corresponding to a percentage error of in the estimated value of $T_{\rm eff}$ ."
 We have optimized ; bv minimizing the RAIS difference in log1 as a function of 3., We have optimized $\beta$ by minimizing the RMS difference in $\log T_{\rm eff}$ as a function of $\beta$.
 In varving 2 from 0 to 1.5. we find that the RAIS clifference goes through a well-defined minimum of ab 3= 0.7. and back up to again: accordingly we select 9=0.7 as the optimum value.," In varying $\beta$ from 0 to 1.5, we find that the RMS difference goes through a well-defined minimum of at $\beta=0.7$ , and back up to again; accordingly we select $\beta=0.7$ as the optimum value."
 The middle panel in Figure 9 shows the result., The middle panel in Figure \ref{fig9} shows the result.
" It is apparent that without the :A, constraint. 6 of the 7 Marshetal.(2010). objects would have been incorrectly assigned temperatures below 2000 Ix. A corresponding plot of our ly estimates versus the previously published values. where available. is shown in Figure 10.."," It is apparent that without the $A_V$ constraint, 6 of the 7 \citet{mar10} objects would have been incorrectly assigned temperatures below 2000 K. A corresponding plot of our $A_V$ estimates versus the previously published values, where available, is shown in Figure \ref{fig10}."
 In order (o assess the sensitivity of our temperature estimation technique to the degree of extinction. we have artificially reddened (he data by various amounts and repeated the estimation procedure.," In order to assess the sensitivity of our temperature estimation technique to the degree of extinction, we have artificially reddened the data by various amounts and repeated the estimation procedure."
 We find that the applied extinction produces very little perturbation in estimated temperature., We find that the applied extinction produces very little perturbation in estimated temperature.
 As an example. the bottom panel of Figure 9 shows the elleet of an applied “ly of 50 mag.," As an example, the bottom panel of Figure \ref{fig9}
 shows the effect of an applied $A_V$ of 50 mag."
 For each estimated temperature. (he models provide a corresponding mass which can be compared with the spectroscopically-estimated mass.," For each estimated temperature, the models provide a corresponding mass which can be compared with the spectroscopically-estimated mass."
" For the voung brown dwarls. the RAIS difference in log mass (log,, M) between our estimates ancl previous spectroscopic determinations was 0.41. corresponding to an average error of a [actor of ~2—3 in mass estimation."," For the young brown dwarfs, the RMS difference in log mass $\log_{10}M$ ) between our estimates and previous spectroscopic determinations was 0.41, corresponding to an average error of a factor of $\sim2-3$ in mass estimation."
" Comparison between (he uncertainties in ζω 4, in Table 1 and (he scatter in these corresponding quantities in Figure 9. indicates that the true uncertaintiesare much larger than theformal errors of maximun likelihood estimation."," Comparison between the uncertainties in $T_{\rm eff}$, $A_V$ in Table 1 and the scatter in these corresponding quantities in Figure \ref{fig9} indicates that the true uncertaintiesare much larger than theformal errors of maximum likelihood estimation."
 The reason is that the latter, The reason is that the latter
In (his case it is important to estimate (he effect of not avoiding the support structure.,In this case it is important to estimate the effect of not avoiding the support structure.
 The level of contamination by (the support vanes for the secondary. mirror can be estimated by deriving the dillraction pattern for a slit with the same proportions of the vanes and using Babinet’s principle., The level of contamination by the support vanes for the secondary mirror can be estimated by deriving the diffraction pattern for a slit with the same proportions of the vanes and using Babinet's principle.
 A vane of width vw along the x-axis and length / in the y-axis has an amplitude of The vane will be brightest along the € axis., A vane of width $w$ along the $x$ -axis and length $l$ in the $y$ -axis has an amplitude of The vane will be brightest along the $\xi$ axis.
 Ideally. support vanes should be rotated. 45° wilh respect to the IICR. so that the diffraction spikes can be masked by the lower contrast regions of the PSF.," Ideally, support vanes should be rotated $^\circ$ with respect to the HCR, so that the diffraction spikes can be masked by the lower contrast regions of the PSF."
" A, can be compared to ACO).", ${\cal A}_s$ can be compared to ${\cal A}(0)$.
" If A,<A(Q). and the secondary is not so lavee that the loss in throughput is great. this method may be preferable."," If ${\cal A}_{s} \ll {\cal A}(0)$, and the secondary is not so large that the loss in throughput is great, this method may be preferable."
 As an example we look at spider vanes that have a width of ~101D. which corresponds to a contrast of 10.7 al a distance of 5A/D assuming (hat the two vanes are oriented 45° (o the axis of interest.," As an example we look at spider vanes that have a width of $\sim10^{-3}D$, which corresponds to a contrast of $10^{-8}$ at a distance of $5\lambda/D$ assuming that the two vanes are oriented $^\circ$ to the axis of interest."
 Clearly (his places a fundamental limit on the width of any support structure (or gaps in a multi-mirror design) for an extrasolar planet search., Clearly this places a fundamental limit on the width of any support structure (or gaps in a multi-mirror design) for an extrasolar planet search.
" Taking the limiting contrast to be 10.!"" al 5A/D. the size limit is 101D."," Taking the limiting contrast to be $10^{-10}$ at $\lambda/D$, the size limit is $10^{-4}D$."
" We designed a mask to be used at the Mt. Wilson 100"" telescope for preliminary observations. as well as a single aperture design for testing in the lab."," We designed a mask to be used at the Mt. Wilson $^{\prime\prime}$ telescope for preliminary observations, as well as a single aperture design for testing in the lab."
 The diameter of the secondary at Mt. Wilson is 7304 the diameter of the primary ancl the width of the spider vanes are ~.25% of the diameter., The diameter of the secondary at Mt. Wilson is $\sim$ the diameter of the primary and the width of the spider vanes are $\sim$ of the diameter.
 We decided to completely avoid the support structure for the initial prototvpe to lower the risk of the PSF being contaminatecl by misalienments of the pupil mask., We decided to completely avoid the support structure for the initial prototype to lower the risk of the PSF being contaminated by misalignments of the pupil mask.
 For a final design we decided to (ry. a variation of what was proposed bv ?.. using a contour based on Equation 13. by placing 3 sub-apertures in each quadrant of the mask to maximize throughput to about," For a final design we decided to try a variation of what was proposed by \citet{spergel01}, using a contour based on Equation \ref{eqn:asym} by placing 3 sub-apertures in each quadrant of the mask to maximize throughput to about."
16%... Table 1 shows the parameters that we used for the two masks as well as the positioning of the apertures for the Mt. Wilson design.," Table \ref{tab:params}
 shows the parameters that we used for the two masks as well as the positioning of the apertures for the Mt. Wilson design."
 Figure 2. shows what the final design looked like., Figure \ref{fig:realmask} shows what the final design looked like.
 Figure 3. shows a comparison between a J band image taken with (he mask and a theoretical PSF modeled by taking digital images of the apertures ab high magnilication and taking a 2D FFT with LDL., Figure \ref{fig:psfcomp} shows a comparison between a J band image taken with the mask and a theoretical PSF modeled by taking digital images of the apertures at high magnification and taking a 2D FFT with IDL.
 Since the spatial scale of the PSF determined in Equation 2 is scalable with wavelength. one can build a multi-wavelength PSF by adding the scaled PSF in wavelength bins together and multipling by the transmission of the particular filter used.," Since the spatial scale of the PSF determined in Equation \ref{eqn:axietared}
 is scalable with wavelength, one can build a multi-wavelength PSF by adding the scaled PSF in wavelength bins together and multiplying by the transmission of the particular filter used."
certain outliers to this relation as given recently by G09 at >3c significance level.,certain outliers to this relation as given recently by G09 at $>3\sigma$ significance level.
" It is notable that the results of our simulations are in contrast with the findings of previous authors, in particular N08 G08, where they find that BATSE trigger threshold, as well as spectral analysis limits on the two planes of have possibly little or no effects on the distributions of BATSE LGRBs on these two planes."," It is notable that the results of our simulations are in contrast with the findings of previous authors, in particular N08 G08, where they find that BATSE trigger threshold, as well as spectral analysis limits on the two planes of have possibly little or no effects on the distributions of BATSE LGRBs on these two planes."
" The reason for the discrepancies should be sought in the values of the limiting parameters that they use in their simulations, such the average ratio of fluence to peak flux of the bursts (FPR) and the nominal durations of the bursts considered therein, all taken from a small fraction ofdetected,analyzed BATSE LGRBs."," The reason for the discrepancies should be sought in the values of the limiting parameters that they use in their simulations, such the average ratio of fluence to peak flux of the bursts (FPR) and the nominal durations of the bursts considered therein, all taken from a small fraction of, BATSE LGRBs."
" In this sense, their analyses suffer from a circular logic problem "," In this sense, their analyses suffer from a circular logic problem \ref{sec:sbep}) )."
"In addition, the use spectral models with fixed photon (82:2).indices in their simulations, results in a severe underestimation of the selection effects "," In addition, the use spectral models with fixed photon indices in their simulations, results in a severe underestimation of the selection effects \ref{sec:pbep}) )."
"The strong evolution of the peak-energy (Ep,ovs) in the (82.3)).light curves of the bursts (e.g. K06, Ryde 1999; Crider et al."," The strong evolution of the peak-energy $\epo$ ) in the light curves of the bursts (e.g. K06, Ryde 1999; Crider et al."
 1999; Band 1997; Crider et al., 1999; Band 1997; Crider et al.
 1997; Liang Kargatis 1996; Ford et al., 1997; Liang Kargatis 1996; Ford et al.
 1995) is another factor that is overlooked in the simulations of G08 N08., 1995) is another factor that is overlooked in the simulations of G08 N08.
 We have worked to make the simulations presented here free from the above mentioned deficiencies., We have worked to make the simulations presented here free from the above mentioned deficiencies.
" Another strong argument that favors an unphysical origin for the Amati, Ghirlanda and possibly other 3-parameter relations, such as the empirical Liang-Zhang relation, comes from inter-comparisons between the proposed relations."," Another strong argument that favors an unphysical origin for the Amati, Ghirlanda and possibly other 3-parameter relations, such as the empirical Liang-Zhang relation, comes from inter-comparisons between the proposed relations."
" Previous authors have reported a significant scatter reduction in transforming the Amati relation to these 3-parameter relations, specifically the Ghirlanda relation."," Previous authors have reported a significant scatter reduction in transforming the Amati relation to these 3-parameter relations, specifically the Ghirlanda relation."
" However, considering the same sample for both relations that are being compared to each other, we find that the scatter reduction and correlation improvements are insignificant (89)."," However, considering the same sample for both relations that are being compared to each other, we find that the scatter reduction and correlation improvements are insignificant \ref{sec:Ghirlanda}) )."
" Therefore, in order to have a meaningful comparison of any two relations with each other, in particular the Amati Ghirlanda relations, it is important to consider the same data set for both relations."," Therefore, in order to have a meaningful comparison of any two relations with each other, in particular the Amati Ghirlanda relations, it is important to consider the same data set for both relations."
" It is also noteworthy that the sample of LGRBs considered by G08 to construct the tight Amati relation also shows a strong correlation in the observer frame, with a scatter comparable to the dispersion in the rest frame Amati relation, differing by only 0.02 dex."," It is also noteworthy that the sample of LGRBs considered by G08 to construct the tight Amati relation also shows a strong correlation in the observer frame, with a scatter comparable to the dispersion in the rest frame Amati relation, differing by only 0.02 dex."
" This indicates that the tightness of the Amati relation is only a ghost of the tight correlation of this sample of LGRBs in the observer frame, reinforced by redshifting the spectral parameters of the bursts from the observer to the rest frame plane."," This indicates that the tightness of the Amati relation is only a ghost of the tight correlation of this sample of LGRBs in the observer frame, reinforced by redshifting the spectral parameters of the bursts from the observer to the rest frame plane."
" For any random redshifts that these bursts might have, the rest frame Amati relation is on average always tighter than the Amati relation in the observer frame (84))."," For any random redshifts that these bursts might have, the rest frame Amati relation is on average always tighter than the Amati relation in the observer frame \ref{sec:discussion}) )."
" Moreover, the apparent frequent inconsistencies of the sub-luminous LGRBs with the Amati relation, appear to have no physical origin and can be attributed purely to redshifting the spectral parameters of the bursts that mainly reside on a narrow strip in the observer frame, by a redshift z<0.2 Figure [7], ϱ))."," Moreover, the apparent frequent inconsistencies of the sub-luminous LGRBs with the Amati relation, appear to have no physical origin and can be attributed purely to redshifting the spectral parameters of the bursts that mainly reside on a narrow strip in the observer frame, by a redshift $z<0.2$ \ref{sec:discussion} Figure \ref{outliers}, \ref{AG08}) )."
" In sum, the Amati relation as proposed by Amati (2002), Amati (2006) Ghirlanda et al. ("," In sum, the Amati relation as proposed by Amati (2002), Amati (2006) Ghirlanda et al. ("
"2008) appears to be greatly affected by complex selection effects in triggering, spectral analyses redshift measurements of LGRBs on the dim side of the plane.","2008) appears to be greatly affected by complex selection effects in triggering, spectral analyses redshift measurements of LGRBs on the dim side of the plane."
 The lack of LGRBs on the soft bright side of the plane might possibly retain an underlying physical origin., The lack of LGRBs on the soft bright side of the plane might possibly retain an underlying physical origin.
" Nevertheless, the practical use of Epos as a standard candle is questioned, as its detector convolutions likely compromise its use as a discerning probe for cosmological models."," Nevertheless, the practical use of $\epo$ as a standard candle is questioned, as its detector convolutions likely compromise its use as a discerning probe for cosmological models."
" This work could not have been accomplished without the vast time and efforts spent by many workers over the past decade, in particular BATSE team, including the designers, builders, and analysts for the burst detectors on board the Gamma, Ray Observatory who have accumulated and analyzed the observations and summarized them in BATSE GRB catalogs."," This work could not have been accomplished without the vast time and efforts spent by many workers over the past decade, in particular BATSE team, including the designers, builders, and analysts for the burst detectors on board the Gamma Ray Observatory who have accumulated and analyzed the observations and summarized them in BATSE GRB catalogs."
" In particular, we acknowledge several useful communications with David Band, and dedicate this paper to his memory."," In particular, we acknowledge several useful communications with David Band, and dedicate this paper to his memory."
galaxies appear to show an absence of such features. with the best example to date being NGC 5548 (Pounds et al.,"galaxies appear to show an absence of such features, with the best example to date being NGC 5548 (Pounds et al."
 2003b)., 2003b).
 Clearly. the presence or absence of a relativistic iron line depends upon currently. unknown factors and is not a simple function of AGN class.," Clearly, the presence or absence of a relativistic iron line depends upon currently unknown factors and is not a simple function of AGN class."
 Progress must be made by careful analysis of as many AGN X-ray spectra às. possible., Progress must be made by careful analysis of as many AGN X-ray spectra as possible.
 With this motivation. thisPaper presents a careful. analysis of the hard-band EPIC-pn spectrum of the Sevlert-l galaxy NGC 4593 (2= 0.009).," With this motivation, this presents a careful analysis of the hard-band EPIC-pn spectrum of the Seyfert-1 galaxy NGC 4593 $z=0.009$ )."
 Section 2 deseribes in. brief our observation and data reduction techniques., Section 2 describes in brief our observation and data reduction techniques.
" Section 3 presents an analysis of the kkeV spectrum of NGC 4593 and. demonstrates. à. marked. absence of a ""stancard relativistic iron line. à result. which is discussed in more detail in Section. 4."," Section 3 presents an analysis of the keV spectrum of NGC 4593 and demonstrates a marked absence of a “standard” relativistic iron line, a result which is discussed in more detail in Section 4."
 Section 5 summarizes our. principal conclusions., Section 5 summarizes our principal conclusions.
 observed NGC 4593 for a total of kksec during orbit 465., observed NGC 4593 for a total of ksec during orbit 465.
 All instruments. were operating during this observation., All instruments were operating during this observation.
 The EPIC-pn was operated in its window mode to prevent photon pile-up. using the medium-thick filter to avoid optical light contamination.," The EPIC-pn was operated in its small-window mode to prevent photon pile-up, using the medium-thick filter to avoid optical light contamination."
 Phe EPIC ALOS-1 camera took data in the fast. uncompressecl timing mode. and the MOS-2 camera operated. in prime partial W2 imaging mode.," The EPIC MOS-1 camera took data in the fast uncompressed timing mode, and the MOS-2 camera operated in prime partial W2 imaging mode."
 Though the MOS results will not. be discussed. further here. they mirrored the pn data within the expected. errors of calibration elfects.," Though the MOS results will not be discussed further here, they mirrored the pn data within the expected errors of calibration effects."
 The average count ratefor⋅ the pn instrument. was 1., The average count ratefor the pn instrument was $^{-1}$.
 The pipeline data were reprocessecl using version 5.4.1 of the Science Analysis Software. and the matching calibration files (CCE from |.2003. January 14. via itpz/xmm.vilspa.ess.es/ccf).," The pipeline data were reprocessed using version 5.4.1 of the Science Analysis Software, and the matching calibration files (CCF from 2003 January 14 via http://xmm.vilspa.esa.es/ccf)."
 From these. we rebuilt the calibration index file usingcifbuild. and then reprocessed he IEPIC data using the ancl tasks.," From these, we rebuilt the calibration index file using, and then reprocessed the EPIC data using the and tasks."
 Dad jxels ancl cosmic rav spikes (narrow time filtering) were removed from the events files via the task within the SAS., Bad pixels and cosmic ray spikes (narrow time filtering) were removed from the events files via the task within the SAS.
 Circular extraction regions aarcsecs in radius were hen defined both on and olf the source for generating source and background spectra. respectively. using the ask.," Circular extraction regions arcsecs in radius were then defined both on and off the source for generating source and background spectra, respectively, using the task."
 Response matrices and ancillary response files were created. using andarfgen. and the data were then grouped with their respective response files using thegrppha such as to impose a minimum of 25in counts per spectral bin.," Response matrices and ancillary response files were created using and, and the data were then grouped with their respective response files using the such as to impose a minimum of 25 counts per spectral bin."
 This binning is required in order to get sullicient counts per bin to make 2 x7-spectral fitting a valid statistical process., This binning is required in order to get sufficient counts per bin to make $\chisq$ -spectral fitting a valid statistical process.
 Spectral modeling and analysis was performed. using the NSPEC package., Spectral modeling and analysis was performed using the XSPEC package.
 In this paper. we analyze only the hard band kkeV) data from the EPIC-pn. camera.," In this paper, we analyze only the hard band keV) data from the EPIC-pn camera."
 This energy restriction prevents us from having to model the soft. excess. ionized absorption. and recombination line emission present in the 2kkeV band.," This energy restriction prevents us from having to model the soft excess, ionized absorption, and recombination line emission present in the keV band."
 The soft spectrum of this object. will be discussed in detail in another publication (Brenneman et al., The soft spectrum of this object will be discussed in detail in another publication (Brenneman et al.
 2004)., 2004).
 Initially. we fit this spectrum with a mocel consisting of a simple power-law modified bv neutral Galactic absorption with a column density of Ny=L97«107em7 (Elvis. Wilkes Lockman 1989).," Initially, we fit this spectrum with a model consisting of a simple power-law modified by neutral Galactic absorption with a column density of $N_{\rm H}=1.97\times 10^{20}\pcmsq$ (Elvis, Wilkes Lockman 1989)."
 The best fitting power-law index is -- (A?fol=1993/1476) giving a model keV (lux and luminosity of £4110Pergem7s ancl 7.601077ergs+ respectively.," The best fitting power-law index is $\Gamma=1.69\pm
0.01$ $\chi^2/{\rm dof}=1993/1476$ ) giving a model keV flux and luminosity of $4.41\times 10^{-11}\ergpcmsqps$ and $7.60\times
10^{42}\ergps$ , respectively."
 This best-fit is shown in lac. The power-law describes the spectrum well apart from two obvious line-like features between kkeV. Modeling these eatures with two Gaussian profiles improves the overall fit significantly (v7/dof=1526/1470)., This best-fit is shown in \ref{fig:spectrum}a a. The power-law describes the spectrum well apart from two obvious line-like features between keV. Modeling these features with two Gaussian profiles improves the overall fit significantly $\chi^2/{\rm dof}=1526/1470$ ).
" Dest-fitting (rest-frame) centroid energies. standard deviations and equivalent wiclths or these two features are: fo)=6.39d0.01keV. σι—1TeV. EW,=LITA1l1eV and E»=6.95+0.05keV. σο=102ot eV. EW.=39418eV."," Best-fitting (rest-frame) centroid energies, standard deviations and equivalent widths for these two features are; $E_1=6.39\pm 0.01\keV$, $\sigma_1=86\pm 17\eV$ , $_1=117\pm 11\eV$ and $E_2=6.95\pm
0.05\keV$, $\sigma_2=102^{+94}_{-79}\eV$ , $_2=39\pm 13\eV$."
" LI seems very likely that these eatures are the Ke emission lines of ""cold"" iron (i.e.. less han Fexvit)) and hydrogen-like iron that are expected. at kkeV. ancl KkeV. respectively."," It seems very likely that these features are the $\alpha$ emission lines of “cold” iron (i.e., less than ) and hydrogen-like iron that are expected at keV and keV, respectively."
 We do not detect a yelium-like iron line at kkeV the upper limit to the equivalent width of any such line is EW<13eV., We do not detect a helium-like iron line at keV — the upper limit to the equivalent width of any such line is $<13\eV$.
 These relatively narrow emission lines likely originate from material that is rather distant from the central black hole., These relatively narrow emission lines likely originate from material that is rather distant from the central black hole.
 The cold iron line is centered on the systemic velocity of NGC 4593 and is well resolved (FWILAI=10900+2200kms 1)., The cold iron line is centered on the systemic velocity of NGC 4593 and is well resolved $10900\pm 2200\kmps$ ).
 For comparison. the broad. optical LL? line in NGC 4593 has EFWIIM-A9104300kms (Cirupe et al.," For comparison, the broad optical $\beta$ line in NGC 4593 has $4910\pm 300\kmps$ (Grupe et al."
 2004). approximately half of the velocity width of the cold iron line.," 2004), approximately half of the velocity width of the cold iron line."
 Thus. it seems clear that the cold iron line is originating from a region that lies significantly inside of the optical broad emission. line region. (ODLIU.. anc hence cannot be identified. with N-rav. reflection from the classic “molecular gas torus” postulated by unified: Sevfer schemes.," Thus, it seems clear that the cold iron line is originating from a region that lies significantly inside of the optical broad emission line region (OBLR), and hence cannot be identified with X-ray reflection from the classic “molecular gas torus” postulated by unified Seyfert schemes."
 We note that the data sets no usefu constraints on the presence of a Compton backscatterec continuum expected from X-ray reflection by cold. matter. thus the possibility remains that this line might be forme bv the X-ray. illumination and subsequent [uorescence. of optically-thin material (as opposed. to the opticallv-thick material normally envisaged in X-ray reflection).," We note that the data sets no useful constraints on the presence of a Compton backscattered continuum expected from X-ray reflection by cold matter, thus the possibility remains that this line might be formed by the X-ray illumination and subsequent fluorescence of optically-thin material (as opposed to the optically-thick material normally envisaged in X-ray reflection)."
 The hyerogen-like iron line. on the other hand. is only ο resolved (EWIINME12200.fikms +).," The hydrogen-like iron line, on the other hand, is only marginally resolved $12200^{+11200}_{-9400}\kmps$ )."
 Thus. it is not possible to conclude. where. relative to the ODLIt. the ionized emission originates.," Thus, it is not possible to conclude where, relative to the OBLR, the ionized emission originates."
e Lt is. however. possible to say something about the physical process uncderlving this emission.," It is, however, possible to say something about the physical process underlying this emission."
 Lone supposes that this line is emitted by collisionally-ionized thermal plasma (described bv the model in NSPEC: Alewe. CGronensehilel van den Oord 1985: Mewe. Lemen van den Oorcl 1986: lxaastra L992: Liedahl. Osterheld Goldstein 1995). the EPIC data. demand that. the. plasma possess," If one supposes that this line is emitted by collisionally-ionized thermal plasma (described by the model in XSPEC; Mewe, Gronenschild van den Oord 1985; Mewe, Lemen van den Oord 1986; Kaastra 1992; Liedahl, Osterheld Goldstein 1995), the EPIC data demand that the plasma possess"
it to tumble to the currently observed. spin-orbit orientation.,it to tumble to the currently observed spin-orbit orientation.
 The pre-SN spin. Sp. the angular monentum produced by the SN ejecta. AS. and the post-SNspint.. Sov. are related by the conservation of angular momentunm To determine AS. we must know Sy aud σον. but we only. know the direction. not the magnitude. of Sy aud the relationship between Sg and the spin measured today is complicated by relativistic precession (Bretonetal.2008).," The pre-SN spin, $\vec{S}_0$, the angular momentum produced by the SN ejecta, $\Delta \vec{S}$, and the post-SN, $\vec{S}_{SN}$, are related by the conservation of angular momentum To determine $\Delta \vec{S}$ , we must know $\vec{S}_0$ and $\vec{S}_{SN}$, but we only know the direction, not the magnitude, of $\vec{S}_0$ and the relationship between $\vec{S}_{SN}$ and the spin measured today is complicated by relativistic precession \citep{Breton2008}."
. However. we cau still place coustraints on AS.," However, we can still place constraints on $\Delta \vec{S}$."
 Relativistic precession causes the individual pulsar spins to precess about the total angular momentuui ol the system. which is approximately parallel to the orbital angular momeutum.," Relativistic precession causes the individual pulsar spins to precess about the total angular momentum of the system, which is approximately parallel to the orbital angular momentum."
 Such. precessiou preserves the augle between the total angular momentuum aud the spin (which is the spin colatitucle). but not the azimuthal orientation.," Such precession preserves the angle between the total angular momentum and the spin (which is the spin colatitude), but not the azimuthal orientation."
 Thus. the colatitude of Ss relative to the normal to the current orbital plane is equal to the colatitude of the current splu—130 degrees.," Thus, the colatitude of $\vec{S}_{SN}$ relative to the normal to the current orbital plane is equal to the colatitude of the current spin—130 degrees."
 Based on tlie spin of pulsar A. the current orbital plane could be tilted at most 11 degrees relative to the pre-SN orbital plane.," Based on the spin of pulsar A, the current orbital plane could be tilted at most $14$ degrees relative to the pre-SN orbital plane."
 Therefore the colatitude of Ss relative to the pre-SN orbital plaue—aucl therefore relative to S8j— at least 116 clegrees., Therefore the colatitude of $\vec{S}_{SN}$ relative to the pre-SN orbital plane—and therefore relative to $\vec{S}_0$ ---is at least 116 degrees.
 This is also the minimum angle between Sj aud AS., This is also the minimum angle between $\vec{S}_0$ and $\vec{\Delta S}$.
 The augular momentuui produced by the SN must be mis-aligned with the progenitor spin., The angular momentum produced by the SN must be mis-aligned with the progenitor spin.
 To date. most SN simulations have focused on non-rotatiug progenitors (forexample.seeBlondiu&Mlezzacappa2007:Rautsiouetal.2011:Wongwatlianarat 2010): it remains to be seen whether the spin produced by the SN from the collapse of a progenitor cau be so significantly. uis-aligned with the progenitors rotation axis.," To date, most SN simulations have focused on non-rotating progenitors \citep[for example,
see][]{Blondin2007,Rantsiou2011,Wongwathanarat2010}; it remains to be seen whether the spin produced by the SN from the collapse of a progenitor can be so significantly mis-aligned with the progenitor's rotation axis."
" . ⊺∐≺↵↕⊽∖⊽↥↽⋯∙⋜↕↥∐↕∩∐↕≺↵∐↕∩⊔∐≺↵↕⋅⊔⋜≹↸↜⊱↥↽∐⋅⋯↕⋅∖↽↥⋟⇂∐⋯≺↵⊽∖⇁⋯∪≪∖∩↥⋜↕∐≺↵⋯↓⋅∩∐⊳∖↕⋜⋃⋅↓⊳∖∪⋅∙⋝≺∟∐∎−∶⋃⊳∖⋯∑≟""A. pa.; . . p> ⋅ ↥⋯↥⊳∖⋜⋃⋅⊟⋅⊳∖⋯≺↵⋜↕⊳∖⋯⋅↩≺⊔∐⋜↕⊳∖⊳∖∩↥∎↕⋅⊇⋅∀∟⊔∙↸⋮⊳∖↩≺↵⊺⋜↕∣≻↥≺↵↽⊔⋜⋯≺⋜↕↥⋅⋜↕≺∐⋃⊳∖∩↥∎∐⊔⊆⋯⋅∐⊳∖∢∙⋃∐⋅≺↵∐↕⊳∖↥↽≻↥∐⋜↕∐∑≟⋃↥⋜⋃⋅ momentum is ⋅⋅∣2x↽LOY5 & emn?2 |."," The typical moment of inertia \citep{Spruit1998} of a neutron star is $0.36MR^2$ ; using pulsar B's measured mass of $1.25 M_{\odot}$ (see Table \ref{tab:system-parameters}) ) and a radius of $10$ km, its current spin angular momentum is $2 \times 10^{45}$ g $^2$ $^{-1}$."
 SinceT this: spin ↽⋅⋅is retrograde whereas the pre-SN.the spin is roughly aligned. (within LI. given pulsar A’s small spin tilt) with the current orbital plane. we can place a lower limit ou the change of angular momenttun neeced to explain pulsar B's large aud retrograde spin tilt. where equality olds when Sj=0.," Since this spin is retrograde whereas the pre-SN spin is roughly aligned (within $^\circ$, given pulsar A's small spin tilt) with the current orbital plane, we can place a lower limit on the change of angular momentum needed to explain pulsar B's large and retrograde spin tilt, where equality holds when $S_0 = 0$."
 Because the angle between Sy auc σον is greater than 90 degrees. aly progenitor spin only increases the amount of angular momentt that must be adde to the pulsar by the kick.," Because the angle between $\vec{S}_0$ and $\vec{S}_{SN}$ is greater than 90 degrees, any progenitor spin only increases the amount of angular momentum that must be added to the pulsar by the kick."
 This is demonstrated geometrically in Figure 2ce. The above discussion bas been fully general., This is demonstrated geometrically in Figure \ref{fig:spin-geometry}c c. The above discussion has been fully general.
 To extract more constraiuts from the observe spin-spin misalignment. we must make some assumptious about the origin of the pulsar spin.," To extract more constraints from the observed spin-spin misalignment, we must make some assumptions about the origin of the pulsar spin."
 As a situplifiecd model to elucidatethe scales involved in this scenario. let us assume that the saine üinpulsive kick linear uomenutum)that changes the orbit of the system is also offset from the," As a simplified model to elucidatethe scales involved in this scenario, let us assume that the same impulsive kick linear momentum)that changes the orbit of the system is also offset from the"
MOST (AMierovariahility Oscillations of οας) is a microsatellite dedicated to detection and characterisation of variability of stars.,MOST (Microvariability Oscillations of STars) is a microsatellite dedicated to detection and characterisation of variability of stars.
 Le uses a 15-cm telescope which feeds a CCD photometer through a single. custom. broadband optical filter (050 700 nm).," It uses a 15-cm telescope which feeds a CCD photometer through a single, custom, broadband optical filter (350 – 700 nm)."
 The pre-launeh characteristics of the mission are described by Walkeretal.(2003). and he initial post-launch performance by Matthews. (2004)., The pre-launch characteristics of the mission are described by \citet{WM2003} and the initial post-launch performance by \citet{M2004}.
. Although the main goal of the mission was to study »ulsations of stars. continuous photometry of selected fields ed to discovery of many new variables of different. tvpes.," Although the main goal of the mission was to study pulsations of stars, continuous photometry of selected fields led to discovery of many new variables of different types."
 Several new variables were found among guide stars used ο accurately point the telescope., Several new variables were found among guide stars used to accurately point the telescope.
 After the failure of the CCD used for the satellite stabilization in January 2006. he number of guide stars utilized for this purpose in the science CCD was increased to a few dozen per field which ed to detection of many variable stars.," After the failure of the CCD used for the satellite stabilization in January 2006, the number of guide stars utilized for this purpose in the science CCD was increased to a few dozen per field which led to detection of many variable stars."
 Observations in a single. wide-band filter ancl small photometric amplitudes often. prevented correct. determination of variability. type or newly discovered. objects.," Observations in a single, wide-band filter and small photometric amplitudes often prevented correct determination of variability type for newly discovered objects."
 The brightness range where spectroscopic information is poor or absent is typically. within NS«V1., The brightness range where spectroscopic information is poor or absent is typically within $8 < V <11$.
 While not. all of some 150 initially selected: targets have been observed in this program. we have decided. to publish our results because the David. Dunlap Observatory was closed on 2008 July 2 and no longer functions.," While not all of some 150 initially selected targets have been observed in this program, we have decided to publish our results because the David Dunlap Observatory was closed on 2008 July 2 and no longer functions."
 Because we have no access to a similar 2-metre class telescope and because the current data form a homogenous anc uniform sample. this survey is published in the present form. as one of the last legacies of this venerable observatory.," Because we have no access to a similar 2-metre class telescope and because the current data form a homogenous and uniform sample, this survey is published in the present form, as one of the last legacies of this venerable observatory."
 The present. paper gives for the survey. targets (1) an estimate of the spectral type. (ii) radial velocity (for most of stars just one observation is available). (iii) the projected rotational velocity. for cases of resin?30 km s.," The present paper gives for the survey targets (i) an estimate of the spectral type, (ii) radial velocity (for most of stars just one observation is available), (iii) the projected rotational velocity, for cases of $v \sin i > 30$ km $^{-1}$."
 The spectral type was estimated on the basis of the best template producing the integral of the broacdenine-funetion (BI) closest to unity (see Section 3))., The spectral type was estimated on the basis of the best template producing the integral of the broadening-function (BF) closest to unity (see Section \ref{analysis}) ).
 For a few stars previously known or discovered. by us to. be spectroscopic binaries. we present radial velocities for individual components ancl preliminary orbits.," For a few stars previously known or discovered by us to be spectroscopic binaries, we present radial velocities for individual components and preliminary orbits."
 MOST. photometry of the newly discovered eclipsing binaries will be presented in a separate paper., MOST photometry of the newly discovered eclipsing binaries will be presented in a separate paper.
 Analvsis of variables (including eclipsing binaries) in the field of M67 is given in Pribullaetal. (2008)., Analysis of variables (including eclipsing binaries) in the field of M67 is given in \citet{m67}. .
. We note that many new variables are ὁ Seuti or 5. Doradus pulsators. where estimates of the spectral type and. the projected rotational velocity. οσαἐν are important data for the asteroscismology mocelling.," We note that many new variables are $\delta$ Scuti or $\gamma$ Doradus pulsators, where estimates of the spectral type and the projected rotational velocity, $v \sin i$, are important data for the asteroseismology modelling."
 Spectroscopic observations were obtained using the slit spectrograph inthe Cassegrain focus of I.58m telescope of, Spectroscopic observations were obtained using the slit spectrograph inthe Cassegrain focus of 1.88m telescope of
paper strongly favour this scenario at least for NGC 7674 and NGC 4968 (see Sect. 3)).,"paper strongly favour this scenario at least for NGC 7674 and NGC 4968 (see Sect. \ref{analysis}) ),"
 another test can be done on the basis of their [OLLI]. infrared (IR) and X-ray fluxes. as explained. for example. by ?..," another test can be done on the basis of their [OIII], infrared (IR) and X-ray fluxes, as explained, for example, by \citet{pb02}."
 Looking at the diagrams shown in Fig. 5..," Looking at the diagrams shown in Fig. \ref{diagrams},"
 it is clear that all the four sources populate or are very close to the regions of the Compton-thick Seyfert galaxies., it is clear that all the four sources populate or are very close to the regions of the Compton-thick Seyfert galaxies.
" However. we note here that it should be better to define them objects. since these diagrams cannot distinguish between highly obscured and ""switehed-off' sources."," However, we note here that it should be better to define them objects, since these diagrams cannot distinguish between highly obscured and `switched-off' sources."
 Indeed. in at least one case. NGC 7674. the diagrams in Fig.," Indeed, in at least one case, NGC 7674, the diagrams in Fig."
" 5. show that the source moved from the Compton-thin to the Compton-thick region between two X-ray observations. thus being a good ""switched off? candidate."," \ref{diagrams} show that the source moved from the Compton-thin to the Compton-thick region between two X-ray observations, thus being a good `switched off' candidate."
 We will discuss this object further in Sect. 4.3.., We will discuss this object further in Sect. \ref{7674off}.
 None of the sources in our sample was caught in a high state. comparable to the one measured by HEAO A-1.," None of the sources in our sample was caught in a high state, comparable to the one measured by HEAO A-1."
 Therefore. before speculating on large flux variations. we should first assess the reliability of the LMA fluxes.," Therefore, before speculating on large flux variations, we should first re-assess the reliability of the LMA fluxes."
 A first test is to check the optical identifications given by ? for the four sources in our sample., A first test is to check the optical identifications given by \citet{grossan92} for the four sources in our sample.
 We reconstructed, We reconstructed
counts. 1t136 in the LMC. shows no evidence [or a non-salpeter slope of the high-niass stars (7).,"counts, R136 in the LMC, shows no evidence for a non-Salpeter slope of the high-mass stars ."
. llowever. his cluster falls below the mass limit. above which the »wametrisationApametruat chosenthosen here.here (eq. 233 3))," However, this cluster falls below the mass limit above which the parametrisation chosen here (eq. \ref{eq:a3M}) )"
" implies implies:a top-heavy,heavy IME.", implies a top-heavy IMF.
 Additionally. no common description is used in the iterature to characterize the top-heaviness.," Additionally, no common description is used in the literature to characterize the top-heaviness."
 Some authors vary the peak of the IAL while others vary. the. slope or the lower and upper mass limit., Some authors vary the peak of the IMF while others vary the slope or the lower and upper mass limit.
 Some available data »onts are shown inFig., Some available data points are shown inFig.
 4. together with the model oedietions., \ref{fig:result} together with the model predictions.
 The cosmologically interpreted observations w dleianglemainBodyCitationEnd3932] απ suggest rather Llat IMES already for relatively low SEIs and are cillieult to reconcile with the here presented. model even for values Of Mean Which are very high (10 ALL).," The cosmologically interpreted observations by and suggest rather flat IMFs already for relatively low SFRs and are difficult to reconcile with the here presented model even for values of $M_\mathrm{ecl,min}$ which are very high $10^5$ $M_\odot$ )."
 But as can also be seen in Fig., But as can also be seen in Fig.
 4 the IGIME results are in good agreement with the constraints for the Galactic and MAI bulee with7). as well as the constrains for the present-cay mass density from. the cosmological SELL fines).," \ref{fig:result} the IGIMF results are in good agreement with the constraints for the Galactic and M31 bulge with, as well as the constrains for the present-day mass density from the cosmological SFH )."
 Vhe GAALA-team finds a very similar trend of decreasing slope with SER as the models in their sample of  40000 galaxies (2)., The GAMA-team finds a very similar trend of decreasing slope with SFR as the models in their sample of $\sim$ 40000 galaxies .
.. Likewise.," Likewise,"
"SO WC approximate Π near resonance as where £2, is the radius of exact resonance.",so we approximate it near resonance as where $R_{\rm r}$ is the radius of exact resonance.
 We thus find Olten we want to know the torque density ddr., We thus find Often we want to know the torque density $\rmd T/\rmd r$.
 For a thin disk with proper surface density X. ie. whose 3-dimensional properdensity is py=X9(z:). the rest. mass per unit radius is where dl represents the volume of a spacelike surface spanning the range from r tor|Ar and m is the unit forware-cireetec normal to this surface.," For a thin disk with proper surface density $\Sigma$ , i.e. whose 3-dimensional proper density is $\rho_0=\Sigma\delta(z)$, the rest mass per unit radius is where $\rmd^3V$ represents the volume of a spacelike 3-surface spanning the range from $r$ to $r+\Delta r$ and ${\bmath n}$ is the unit forward-directed normal to this surface."
" E""king the surface to be at constant {. the normal is p,=(07.0.0.0) and the volume element is dV—e""""drdoócdz."," Taking the surface to be at constant $t$, the normal is $n_\alpha = (-\rme^\nu, 0,0,0)$ and the volume element is $\rmd^3V=\rme^{\psi+\mmu}\,\rmd r\,\rmd\phi\,\rmd z$."
 Vhis leads to the result Lt follows that Having the formal solution for the torque (I5q. 501) , This leads to the result It follows that Having the formal solution for the torque (Eq. \ref{eq:TX1}) )
is onlv part of the problem: we also need the resonant. aniplitucle Φον, is only part of the problem; we also need the resonant amplitude ${\cal S}^{(m)}$.
 Ες section evaluates the amplitude and then shows that (within some restrictions)it is gauge-invariant., This section evaluates the amplitude and then shows that (within some restrictions)it is gauge-invariant.
 Llere we require both the perturbation Hamiltonian. and its derivatives with respect to r and νε.," Here we require both the perturbation Hamiltonian, and its derivatives with respect to $r$ and $p_r$."
 These are all to be evaluated at the unperturbed circular orbit using Eq. (46))., These are all to be evaluated at the unperturbed circular orbit using Eq. \ref{eq:H131}) ).
" For Lf, itself: we see that since fy=p&CR) and p,= HEMETAR For the partial derivatives with respect to the Coordinates. we find that in general where the last termi is associated with the cerivative of the denominator in Eq. )."," For $H_1$ itself, we see that since $H_0=\mass{\cal E}(R)$ and $p_\phi=\mass{\cal L}(R)$ , For the partial derivatives with respect to the coordinates, we find that in general where the last term is associated with the derivative of the denominator in Eq. \ref{eq:H131}) )."
" Since OyfOr=O on a circular orbit. this implies For the partial derivatives with respect to the momenta. the general expression is For the specific case of p. we note that at the circular orbit. OHo/O0p,=0 and g=0. so We may now assemble the pieces to compute 5nd, Note that this is independent of the particle mass fr ancl linear in the perturbation A"," Since $\partial H_0/\partial r =0$ on a circular orbit, this implies For the partial derivatives with respect to the momenta, the general expression is For the specific case of $p_r$, we note that at the circular orbit $\partial H_0/\partial p_r=0$ and $g^{tr}=0$, so We may now assemble the pieces to compute ${\cal S}^{(m)}$: Note that this is independent of the particle mass $\mass$ and linear in the perturbation $h^{\alpha\beta}$."
 It is possible to rewrite Eq. (98)), It is possible to rewrite Eq. \ref{eq:SM-explicit}) )
 in terms of the civeular ≟−∖⇁∢⋅⇂⋯∼∐∙∖⇁∐↿⊳↳∖↓⊔↓↿↓, in terms of the circular 4-velocity ${\bmath w}$.
↓≻↓∙∖⇁↓⊔⋏∙≟⇂↓∐⋅∪⊔⋏∙≟⇂↥∣⋡∙∖⇁⊓⊽∶∢⋅∣↿∖↼↳∣∣∽↙⋎∠∕⊐. ⋠⋠ ∕⊐∣↴∖⇁ elves This form will be most. useful. in »oving5 5gaugeὃν invariance and in practical applications., Multiplying through by $w^t = \rme^{-2\nu}({\cal E}-\oomega{\cal L})$ gives This form will be most useful in proving gauge invariance and in practical applications.
" In general the perturbation f°"" could be expressed in many choices of gauge.", In general the perturbation $h^{\alpha\beta}$ could be expressed in many choices of gauge.
 These diller bv the relation Since Eq. (99)), These differ by the relation Since Eq. \ref{eq:FSM}) )
" is linear in f,s. the contributions to SU. from the pre-existing and. gauge perturbations simply add. so to show invariance of the torque it is sullicient to prove that a pure gauge perturbation leads to zero resonant amplitude: S7."," is linear in $h_{\alpha\beta}$, the contributions to ${\cal S}^{(m)}$ from the pre-existing and gauge perturbations simply add, so to show invariance of the torque it is sufficient to prove that a pure gauge perturbation leads to zero resonant amplitude ${\cal S}^{(m)}$."
 We restrict our attention to gauges that preserve the fundamental svnunetrics of the problem. ic. that have rellection across the equatorial plane and have helical symmetry. where the Fourier component has an oscillatory time dependence xePS ," We restrict our attention to gauges that preserve the fundamental symmetries of the problem, i.e. that have reflection across the equatorial plane and have helical symmetry, where the $m$ Fourier component has an oscillatory time dependence $\propto \rme^{-\rmi \Omega_{\rm s}mt}$ ."
AVithout loss of generality. we may consider the Fouricr modes one at a time. so we will consider the order m Fourier mode below andavoid writing the superscript nri ," Without loss of generality, we may consider the Fourier modes one at a time, so we will consider the order $m$ Fourier mode below andavoid writing the superscript $^{(m)}$ "
using the 2.8 keV band flux from the uuresolved cluission in regions οἱ aud B.,using the $2-8$ keV band flux from the unresolved emission in regions $A$ and $B$.
" The derived total hard baud unresolvedN-ray luminosity of Ly=3.9«1019orest corresponds to Ly/=1.0.107""erestLi.."," The derived total hard band unresolvedX-ray luminosity of $L_X=3.9 \times 10^{40} \ \rm{erg \ s^{-1}}$ corresponds to $L_X/L_K=1.0\times10^{29} \ \rm{erg \ s^{-1} \ L_{K,\odot}^{-1} }$."
 Below the 7.241075erest sensitivity limit. the average LAINB huninosity function (Calfanov2001). predicts Ly/Ljy=5.0«107cresPL! du the 2S keV energy baud. assubune an average power law LAINB spectrui with slope of P=1.56 and Galactic cola deusitv. (winetal. 2003)...," Below the $7.2 \times10^{38} \ \rm{erg \ s^{-1}}$ sensitivity limit, the average LMXB luminosity function \citep{gilfanov04} predicts $L_X/L_K=5.9\times10^{28} \ \rm{erg \ s^{-1} \ L_{K,\odot}^{-1} }$ in the $2-8$ keV energy band, assuming an average power law LMXB spectrum with slope of $\Gamma=1.56$ and Galactic column density \citep{irwin03}. ."
 Thus. the Iuniunositv fuuction. of LAINBs has to be scaled up by z70%.," Thus, the luminosity function of LMXBs has to be scaled up by $\approx70\%$."
 Therefore the total predicted umber of LAINBs is 2.2 in regions A aud D., Therefore the total predicted number of LMXBs is $2.2$ in regions $A$ and $B$.
 Taking this as the expectation valuc. the total uuuber of predicted LAINBs and CNXBs is 6.3. whereas 20 sourcesare observed.," Taking this as the expectation value, the total number of predicted LMXBs and CXBs is $6.3$, whereas $20$ sourcesare observed."
 Au additional major uncertainty in the average LAINB Iuimuinositv function is the relatively low umber of bright (Ly>8<LOere& +) sources in the sample of Cilfanov (2001)., An additional major uncertainty in the average LMXB luminosity function is the relatively low number of bright $L_X \gtrsim8 \times10^{38} \ \rm{erg \ s^{-1}} $ ) sources in the sample of \citet{gilfanov04}.
.. This imiplies large systematic uncertainties at the bright cud of the Iuuunosity function., This implies large systematic uncertainties at the bright end of the luminosity function.
 Indeed. according to Cilfanov(2001). the number of bireht LAINBs above Ly>=8οeres| anay he factor of ~| times higher within 90% confidence interval due to svstematie errors.," Indeed, according to \citet{gilfanov04} the number of birght LMXBs above $L_X\gtrsim8 \times10^{38} \ \rm{erg \ s^{-1}} $ may be factor of $\sim4$ times higher within $90\%$ confidence interval due to systematic errors."
 Taking this uncertainty mto account. it is feasible that sienificautly more. altogether 8— 9. resolved sources are bright LAINBsin regions ; aud B.," Taking this uncertainty into account, it is feasible that significantly more, altogether $8-9$ , resolved sources are bright LMXBsin regions $A$ and $B$ ."
 Therefore LAINBs and CNB sources iav account for up to 11 sources of the 20 detected., Therefore LMXBs and CXB sources may account for up to $11$ sources of the $20$ detected.
test for the nature ofLL? broadening.,test for the nature of$\beta$ broadening.
 Systematically larger 1111 masses in ANL systems would then indicate that 117 is due to non-virial motion such as an outllow. while continued agreement would indicate that remains virial. perhaps broadened because ANL systems produce. the proper ionization potential for H2 at a different radius.," Systematically larger $\beta$ masses in ANL systems would then indicate that $\beta$ is due to non-virial motion such as an outflow, while continued agreement would indicate that $\beta$ remains virial, perhaps broadened because ANL systems produce the proper ionization potential for $\beta$ at a different radius."
 The ? catalog produces Mgll-based. virial masses using the prescription of 7. (along other estimators). which has been previously shown to agree with L2 masses to within a scatter of ~0.22 dex (2)..," The \citet{Shen2011} catalog produces -based virial masses using the prescription of \citet{McLure2004} (along other estimators), which has been previously shown to agree with $\beta$ masses to within a scatter of $\sim 0.22$ dex \citep{Shen2008}."
 However. these catalogs restrict not just OL) but also the narrow component of 113 to :1200 km/s in FWIIM.," However, these catalogs restrict not just ] but also the narrow component of $\beta$ to $\leq 1200$ km/s in FWHM."
 ANLs include a broader narrow component. so restricting the ΑΛΛΗΣΤΕ to 1200 Kms will also unelerestimate the narrow component width.," ANLs include a broader $\beta$ narrow component, so restricting the FWHM to 1200 km/s will also underestimate the $\beta$ narrow component width."
 As à result. the line Hux assumed by the fit to represent the broad. component would actually have been a combination of the broad component ancl the wines of the narrow component. resulting in a broad component fit that was artificially too narrow.," As a result, the line flux assumed by the fit to represent the $\beta$ broad component would actually have been a combination of the broad component and the wings of the narrow component, resulting in a broad component fit that was artificially too narrow."
 An underestimate of the LL? EAWHIAL will result in an underestimate in the virial mass Meux EWIIM., An underestimate of the $\beta$ FWHM will result in an underestimate in the virial mass $M_{BH} \propto \textrm{FWHM}^2$ .
 Using the fitting routine described in 2.. we produce an alternative set of masses from the broad LL? component. where the narrow component is no longer constrained to lie at à EWIIM lessthan 1200 km/s. For narrow ΟΠΗ] lines. there is good. agreement between both sets of 11.1 masses. but for ANLs. the ? masses are systematically lower (Fig. 3)).," Using the fitting routine described in \ref{sec:sample}, we produce an alternative set of masses from the broad $\beta$ component, where the narrow component is no longer constrained to lie at a FWHM lessthan 1200 km/s. For narrow ] lines, there is good agreement between both sets of $\beta$ masses, but for ANLs, the \citet{Shen2011} masses are systematically lower (Fig. \ref{fig:hbcomp}) )."
 This is likely due to the ellect described above. where a 7 fit a combination of the broad. anc narrow component as their. broad. component.," This is likely due to the effect described above, where a \citet{Shen2011} fit a combination of the broad and narrow component as their broad component."
 Objects are only included: in iis sample if they have high enough signal-to-noise for our routine to produce a high-quality fit., Objects are only included in this sample if they have high enough signal-to-noise for our routine to produce a high-quality fit.
 Note that the high signal-to-noise sample contains a larger NL fraction than 16 overall sample. since ANLs are on average more Luminous jan other quasars.," Note that the high signal-to-noise sample contains a larger ANL fraction than the overall sample, since ANLs are on average more luminous than other quasars."
 For 4716 high signal-to-noise ANLs at Q4<zES (where both LL2 andAlell masses are available). our 1L2 masses average 0.474 dex higher than ?/ using 11. while for 5417 other quasars our masses average 0.075 dex ugher.," For 4716 high signal-to-noise ANLs at $0.4 < z < 0.8$ (where both $\beta$ and masses are available), our $\beta$ masses average $0.474$ dex higher than \citet{Shen2011} using $\beta$, while for 5417 other quasars our masses average $0.075$ dex higher."
 Some of this cliscrepaney is due to changes in the template used and other line-fitting details. which are xwticularlv important forMegll.," Some of this discrepancy is due to changes in the template used and other line-fitting details, which are particularly important for."
 Compared to the ο Mgll masses. our masses average 0.455 dex higher for ANLs and 170 dex higherfor other quasars.," Compared to the \citet{Shen2011} masses, our masses average $0.455$ dex higher for ANLs and $0.170$ dex higherfor other quasars."
 Since the mass ancl luminosity distributions change with recshift (2).. it is likely best to compare these. mass," Since the mass and luminosity distributions change with redshift \citep{Steinhardt2010a}, , it is likely best to compare these mass"
is roughly along the direction of the large scale filament feeding the cluster.,is roughly along the direction of the large scale filament feeding the cluster.
 A third less massive structure is located towards the west., A third less massive structure is located towards the west.
 Cold dark matter hierarchical structure formation models (e.g. Colberg et al 1999) predict that clusters of galaxies grow via group aceretion., Cold dark matter hierarchical structure formation models (e.g. Colberg et al 1999) predict that clusters of galaxies grow via group accretion.
 In this context. the cluster substructures detected along the path of the large scale filament are probably recent infallen groups.," In this context, the cluster substructures detected along the path of the large scale filament are probably recent infallen groups."
 We also quantified the spatial distribution of galaxies with photo-z«0.2 as a function of radius both for the fLSBs and for the entire sample of galaxies in our images (including non-fLSB galaxies).," We also quantified the spatial distribution of galaxies with $-z
< 0.2$ as a function of radius both for the fLSBs and for the entire sample of galaxies in our images (including non-fLSB galaxies)."
 For this. we counted the numbers of galaxies in concentric annuli with radit varying between 5 and 30 aremin in steps of5 aremin.," For this, we counted the numbers of galaxies in concentric annuli with radii varying between 5 and 30 arcmin in steps of 5 arcmin."
 The density distributions (number of galaxies per aremin?) thus obtained are drawn in Fig. 6.., The density distributions (number of galaxies per $^2$ ) thus obtained are drawn in Fig. \ref{fig:densproflsbs}.
 These distributions were fit by a King model and we added a constant background to take into account non-cluster photo-z«0.2 galaxies: For the 142 fLSBs with photo-z«0.2. the best fit was obtained for the following parameters: P(0)=(6.846.3) 107 arcmin7. PCI)20.1250.03 aremin7. P(2)211.19x0.03 aremin.," These distributions were fit by a King model and we added a constant background to take into account non-cluster $-z<0.2$ galaxies: For the 142 fLSBs with $-z<0.2$, the best fit was obtained for the following parameters: $\pm$ 6.3) $^{-3}$ $^{-2}$, $\pm$ 0.03$^{-2}$ , $\pm$ 0.03 arcmin."
 For all the galaxies. the corresponding numbers are: P(0)21.8720.09. aremin. 1)=2.4540.54 ürcmin7. P(2)25.29£0.5] aremin.," For all the galaxies, the corresponding numbers are: $\pm 0.09$ $^{-2}$, $\pm 0.54$ $^{-2}$, $\pm 0.51$ arcmin."
 As can be seen in Fig. 6..," As can be seen in Fig. \ref{fig:densproflsbs},"
 fLSBs in Abell 496 are not uniformly distributed. but are preferentially found toward the cluster center. except for a decrease of the number of fLSBs per aremin?. in the cluster innermost point.," fLSBs in Abell 496 are not uniformly distributed, but are preferentially found toward the cluster center, except for a decrease of the number of fLSBs per $^2$ in the cluster innermost point."
 This data point is most likely low because it is located inside the central galaxy radius. thus fLSBs in this region would have been missed by our analysis.," This data point is most likely low because it is located inside the central galaxy radius, thus fLSBs in this region would have been missed by our analysis."
 Note that although the fLSBS are concentrated toward the cluster center they are less concentrated than the whole galaxy population. as discussed in Section 4..," Note that although the fLSBS are concentrated toward the cluster center they are less concentrated than the whole galaxy population, as discussed in Section \ref{sec:discu}."
 Before computing a luminosity function. for fLSBs. it is important to investigate the completeness level of our sample.," Before computing a luminosity function for fLSBs, it is important to investigate the completeness level of our sample."
 We show in Fig., We show in Fig.
 7. the magnitude histogram of the 783 fLSBs with available photo-zs and the luminosity function of the 142 fLSBs with photo-z«0.2., \ref{fig:histos} the magnitude histogram of the 783 fLSBs with available $-z$ s and the luminosity function of the 142 fLSBs with $-z<0.2$.
 The peak of the magnitude histogram for the 783 fLSBs ts located close to +’=22.7., The peak of the magnitude histogram for the 783 fLSBs is located close to $r'=22.7$.
 This gives a first estimate of the completeness limit of the sample., This gives a first estimate of the completeness limit of the sample.
 We also performed simulations in order to have an independent estimate of the completeness in the » images., We also performed simulations in order to have an independent estimate of the completeness in the $r'$ images.
 The simulation adds artificial objects of different shapes and magnitudes to the CCD images and then attempts to recover them by running SExtractor again with the same parameters used for primary object detection (see Adami et al., The simulation adds artificial objects of different shapes and magnitudes to the CCD images and then attempts to recover them by running SExtractor again with the same parameters used for primary object detection (see Adami et al.
 2006 for more details)., 2006 for more details).
 In this way. the completeness is measured on the original CCD frames.," In this way, the completeness is measured on the original CCD frames."
 We estimated the completeness of our catalog for fLSBs using simulated point-like objects with a Gaussian profile of FWHM 3.3 aresec (o-1.4 aresec)., We estimated the completeness of our catalog for fLSBs using simulated point-like objects with a Gaussian profile of FWHM 3.3 arcsec $\sigma$ =1.4 arcsec).
 This is the typical maximal size of a fLSB in our catalog (see Fig. 8))., This is the typical maximal size of a fLSB in our catalog (see Fig. \ref{fig:kr}) ).
 We also divided the full field of view in 100 different sub-regions to have the completeness at different locations 1n the cluster., We also divided the full field of view in 100 different sub-regions to have the completeness at different locations in the cluster.
 The percentage of recovered fLSBs as a function of the 7 magnitude is shown in Fig. 9..," The percentage of recovered fLSBs as a function of the $r'$ magnitude is shown in Fig. \ref{fig:simu},"
 where error bars show the variation among these 100 regions., where error bars show the variation among these 100 regions.
" We can see that we reach a completeness at 7? ~22.8,", We can see that we reach a completeness at $r'\sim$ 22.8.
) This estimate ts similar to the value of the peak of the fLSB magnitude histogram., This estimate is similar to the value of the peak of the fLSB magnitude histogram.
 It also represents an underestimate of the true fLSB completeness level. as most f£LSBs are more compact than a 3.3 aresec FWHM Gaussian profile. and therefore easier to detect.," It also represents an underestimate of the true fLSB completeness level, as most fLSBs are more compact than a 3.3 arcsec FWHM Gaussian profile, and therefore easier to detect."
 We can see in Fig., We can see in Fig.
 7 that the luminosity function decreases for magnitudes fainter than the completeness limit. as expected.," \ref{fig:histos} that the luminosity function decreases for magnitudes fainter than the completeness limit, as expected."
 A power-law fit of the bright part of the lummosity function gives a mean slope of —1.2x0.1., A power-law fit of the bright part of the luminosity function gives a mean slope of $-1.2\pm 0.1$.
 This is significantly shallower than the global luminosity function of Boué et al. (, This is significantly shallower than the global luminosity function of Boué et al. (
2008). who found a slope of —1.55+ 0.05.,"2008), who found a slope of $-1.55 \pm 0.05$ ."
 This means thatthe fLSBs cannot be responsible for all the increase of the global galaxy luminosity function. of the cluster at faint magnitudes., This means thatthe fLSBs cannot be responsible for all the increase of the global galaxy luminosity function of the cluster at faint magnitudes.
 We, We
"For point sources, these observations will be unbiased and the data from all objects in the sky, bellow a certain limiting magnitude, will be sent to the ground.","For point sources, these observations will be unbiased and the data from all objects in the sky, bellow a certain limiting magnitude, will be sent to the ground."
" Certainly, among all those objects, not only galactic sources will be present, but also extragalactic ones."," Certainly, among all those objects, not only galactic sources will be present, but also extragalactic ones."
" In particular, it is expected that point sources up to magnitude 20, in the Gaia passbandG®,, will be “windowed” and"," In particular, it is expected that point sources up to magnitude 20, in the Gaia passband, will be “windowed” and."
"transferred.. As seen in Fig. 2,,"," As seen in Fig. \ref{fig:GRBafterglow},"
 some of the OA events are expected to remain above this limiting magnitude for a certain amount of time., some of the OA events are expected to remain above this limiting magnitude for a certain amount of time.
 The question that remains is whether their duration (at G<20) is enough for them to be observed at a reasonable rate., The question that remains is whether their duration (at $\leqslant$ 20) is enough for them to be observed at a reasonable rate.
" Only two quantities play an important role in estimating the probability of Gaia observing single event from a Pop IIL2: the time whichen the OA remains brighter than G=20, At, and the coordinates (ἶραι,bgat) where the event takes place in the sky."," Only two quantities play an important role in estimating the probability of Gaia observing single event from a Pop III.2: the time whichen the OA remains brighter than G=20, $\Delta t$, and the coordinates $(l_{gal}, b_{gal})$ where the event takes place in the sky."
" Since those quantities are continuous distributions, it is necessary to analyze how the observation probability depends on them by building P(At,Iga,bgat)."," Since those quantities are continuous distributions, it is necessary to analyze how the observation probability depends on them by building $P(\Delta t, l_{gal}, b_{gal})$."
" In the present work, we proceed as follows."," In the present work, we proceed as follows."
" For a given coordinate in the sky, we start by computing the inverse Gaia scanning law to derive a transit time list comprising the instants when Gaia’s telescopes will be pointing at that coordinate."," For a given coordinate in the sky, we start by computing the inverse Gaia scanning law to derive a transit time list comprising the instants when Gaia's telescopes will be pointing at that coordinate."
" To be as realistic as possible, we adopt the Gaia Data Processing and Analysis Consortium’s nominal implementation of it."," To be as realistic as possible, we adopt the Gaia Data Processing and Analysis Consortium's nominal implementation of it."
" Then, we randomly select a point in time during the entire mission lifetime in order to place an event of a certain duration At."," Then, we randomly select a point in time during the entire mission lifetime in order to place an event of a certain duration $\Delta t$."
" Using the transit time list, we check if that event was observed, considering a time window of 4.4 seconds around each transit; this is the time needed for the signal to cross the detection CCD and enter the confirmation CCD."," Using the transit time list, we check if that event was observed, considering a time window of 4.4 seconds around each transit; this is the time needed for the signal to cross the detection CCD and enter the confirmation CCD."
" If there is a superposition between the event duration and this time window, the event is considered detected."," If there is a superposition between the event duration and this time window, the event is considered detected."
" This procedure is then repeated until the estimation of the detection probability, which is derived by simply dividing the number of detected events by the total, does not vary more than between iterations."," This procedure is then repeated until the estimation of the detection probability, which is derived by simply dividing the number of detected events by the total, does not vary more than between iterations."
" Finally, the whole procedure is repeated for each event duration At."," Finally, the whole procedure is repeated for each event duration $\Delta t$."
" As a consequence, we obtain an adequate time-sampling of the P(At,l541,bgat) distribution."," As a consequence, we obtain an adequate time-sampling of the $P(\Delta t, l_{gal}, b_{gal})$ distribution."
" For the determination of the number of OA events observed by Gaia on the entire sky, the coordinate dependency can be averaged out, allowing P(At,lgu,bgat)~P(At)+ε."," For the determination of the number of OA events observed by Gaia on the entire sky, the coordinate dependency can be averaged out, allowing $P(\Delta t, l_{gal}, b_{gal}) \sim P(\Delta t) \pm \epsilon$."
" This is possible because the scanning law is mostly known, so we can reasonably assume that the OA events take place randomly in the sphere."," This is possible because the scanning law is mostly known, so we can reasonably assume that the OA events take place randomly in the sphere."
" The procedure described above was repeated for several positions on the sphere, and the mean and the standard deviation at each event duration were computed."," The procedure described above was repeated for several positions on the sphere, and the mean and the standard deviation at each event duration were computed."
" To allow a good spatial sampling for the estimation of P(At)+e, we tessellate the celestial sphere at the hierarchical triangular mesh, level 4 (Kunsztetal.2001).."," To allow a good spatial sampling for the estimation of $P(\Delta t) \pm \epsilon$, we tessellate the celestial sphere at the hierarchical triangular mesh, level 4 \citep{Kunszt:2001p8829}."
 This means that the simulations were performed at the center of 2048 triangles of approximately equal areas., This means that the simulations were performed at the center of 2048 triangles of approximately equal areas.
" Finally, to obtain the probabilities for the whole sky, an additional effect must be taken into account: the structure of our own Galaxy."," Finally, to obtain the probabilities for the whole sky, an additional effect must be taken into account: the structure of our own Galaxy."
" Since the OAs are extragalactic events, the probability of observation at the galactic plane or bulge should be null or very small, due to extinction and crowding."," Since the OAs are extragalactic events, the probability of observation at the galactic plane or bulge should be null or very small, due to extinction and crowding."
" In this work, we conservatively assumed a null value for the probability of OAsbeing observed at such regions of the sky (defined here as |b|<15° for 345?«|€15? and |b|€5° otherwise)."," In this work, we conservatively assumed a null value for the probability of OAsbeing observed at such regions of the sky (defined here as $|b|\le15^\circ$ for $345^\circ\le l \le 15^\circ$ and $|b|\le5^\circ$ otherwise)."
" The final results, representing the behavior of P(At)+te can be seen in Fig. 8.."," The final results, representing the behavior of $P(\Delta t) \pm \epsilon$ can be seen in Fig. \ref{fig:sphere}. ."
 In accordance with upper limit showed in Fig., In accordance with upper limit showed in Fig.
" 1 and results from deSouzaetal.(2011), we expect between —10?5x events per year in all the sky."," \ref{fig:SFRII} and results from \citet{rafael2011}, we expect between $\sim 10^2-5\times10^3$ events per year in all the sky."
 The uncertainties come from our poor understanding about the efficiency with which gas is converted into stars and GRBsare triggered (two unknown factors for Pop III stars)., The uncertainties come from our poor understanding about the efficiency with which gas is converted into stars and GRBsare triggered (two unknown factors for Pop III stars).
" For a good statistics, we create a mock sample of 105 events randomly generated by the Monte Carlo method in order to infer the PDF of an event to stay below G—20 over At(days)."," For a good statistics, we create a mock sample of $10^5$ events randomly generated by the Monte Carlo method in order to infer the PDF of an event to stay below $G = 20$ over $\Delta t (\rm days)$."
" The average behavior for on-axis and off-axis afterglows as a function of E;,, distribution is shown in Figs. 9--10..", The average behavior for on-axis and off-axis afterglows as a function of $E_{iso}$ distribution is shown in Figs. \ref{fig:probdeltat54}- \ref{fig:probdeltat55}. .
" Once we have P(At), we can generate a sample with 104 events several times and test against their"," Once we have $P (\Delta t)$, we can generate a sample with $10^4$ events several times and test against their"
There are currently three favored models for the formation of GRBs: mergers of ueutron star binaries (Narayan. Paczyuski aud Piran 1992). collapsars (or lypernovae. Woosley 1993. Paczyuski 1998). aud SupraNovae (Vietri and Stella 1998. 1999). but a clearcut test to which each of these could be subjected is still missing.,"There are currently three favored models for the formation of GRBs: mergers of neutron star binaries (Narayan, Paczynski and Piran 1992), collapsars (or hypernovae, Woosley 1993, Paczynski 1998), and SupraNovae (Vietri and Stella 1998, 1999), but a clearcut test to which each of these could be subjected is still missing."
 This is so because the details of the fireball model are to a large extent. iudepeudeut of the details of the euergy deposition mechanism., This is so because the details of the fireball model are to a large extent independent of the details of the energy deposition mechanism.
 The situation has also a j»aradoxical side: even if we managed to observe tlie event that deposits the energy. the euormous expected optical depths (exceeding LOM) would prevent us from deriviug any information on the iiture of the source.," The situation has also a paradoxical side: even if we managed to observe the event that deposits the energy, the enormous expected optical depths (exceeding $10^{10}$ ) would prevent us from deriving any information on the nature of the source."
 Gravitational waves cau probably provide a siguature of at least some of he scenarios. but detection of these from these distant. sources does not look like an imunediate orospect.," Gravitational waves can probably provide a signature of at least some of the scenarios, but detection of these from these distant sources does not look like an immediate prospect."
 For these reasons. παν authors (Botttcher 451999. Mésszárros aud Rees 19983) have proposed hat interaction of the burst with surroundingOm material mighte provide detectable signaturese of star-forming environments.," For these reasons, many authors (Bötttcher 1999, Mésszárros and Rees 1998) have proposed that interaction of the burst with surrounding material might provide detectable signatures of star–forming environments."
 However. the possible discovery. of iron e1uission lines (Piroal... 1999. Yoshidaaf... 1999) from two bursts has highlighted the weak side of this possibility.£e... that any conclusion is subject to the usual uncertainties associated with mocleling[n] complex and incomplete sets of observatious (Lazzatial... 1999. VietriaL... 1999).," However, the possible discovery of iron emission lines (Piro, 1999, Yoshida, 1999) from two bursts has highlighted the weak side of this possibility, that any conclusion is subject to the usual uncertainties associated with modeling complex and incomplete sets of observations (Lazzati, 1999, Vietri, 1999)."
 E show in this paper that a surprisingly simple and robust test exists. for both the neutron binary merger and the collapsar scenarios. resulting [rou a trivial. but so [ar overlooked property of fireballs.," I show in this paper that a surprisingly simple and robust test exists, for both the neutron binary merger and the collapsar scenarios, resulting from a trivial, but so far overlooked property of fireballs."
ο erg em? s7!.,$^{-12}$ erg $^{-2}$ $^{-1}$.
 We derived a luminosity of 110°° ere s7!. or 400Ls.. for NG6A at H(J).," We derived a luminosity of $^{36}$ erg $^{-1}$, or 400, for N66A at $\beta$ )."
 This luminosity corresponds to a flux of 110% photons $ ! ora Lyman continuum flux of .11075 photons s7! for the star. assuming that the rregion is tonization bounded.," This luminosity corresponds to a flux of $^{47}$ photons $^{-1}$ , or a Lyman continuum flux of $^{48}$ photons $^{-1}$ for the star, assuming that the region is ionization bounded."
 The exciting star needed to provide this flux is of spectral type about VV etal..2005) (see also Sect., The exciting star needed to provide this flux is of spectral type about V \citep[][]{Martins05} (see also Sect.
 3.3)., 3.3).
 A number of the derived physical parameters of N66A are summarized in reftab:param.., A number of the derived physical parameters of N66A are summarized in \\ref{tab:param}.
 The mean angular radius of the rregion. corresponding to the FWHM of cross-cuts through the limage. is given in Col.," The mean angular radius of the region, corresponding to the FWHM of cross-cuts through the image, is given in Col."
 1., 1.
 The corresponding linear size. obtained using a distance modulus ofm-M = 18.94 mag (Laney&Stobie.1994) is presented in Col.," The corresponding linear size, obtained using a distance modulus of = 18.94 mag \citep[][]{Laney94} is presented in Col."
 2.The dereddened fflux obtained from a reddening coefficient .) = 0.36 is given in Col., 2.The dereddened flux obtained from a reddening coefficient ) = 0.26 is given in Col.
 3., 3.
 The electron temperature calculated from the forbidden lines ratio 44363/(4959 + 5007). with an accuracy of4%.. is given in Col.," The electron temperature calculated from the forbidden lines ratio 4363/(4959 + 5007), with an accuracy of, is given in Col."
 4., 4.
 The electron density. estimated from the ratio of the ddoublet .66717/6731. 1s presented in Col.," The electron density, estimated from the ratio of the doublet 6717/6731, is presented in Col."
 5., 5.
 It is accurate to880%., It is accurate to.
.. It is well-known that the llines characterize the low-density peripheral zones of rregions (see below for corroboration)., It is well-known that the lines characterize the low-density peripheral zones of regions (see below for corroboration).
 Column 6 gives the rms electron density calculated from the fflux. the radius. r. and the electron temperature.Το. assuming that the rregion is an ionization-bounded Str6mmeren sphere.," Column 6 gives the rms electron density calculated from the flux, the radius, $r$, and the electron temperature, assuming that the region is an ionization-bounded Strömmgren sphere."
 For comparison reasons. an electron density of + 330 em7* was found from our long-slit spectra far from N66À. towards the inner parts of the N66 giant region.," For comparison reasons, an electron density of $\pm$ 30 $^{-3}$ was found from our long-slit spectra far from N66A, towards the inner parts of the N66 giant region."
 This agrees well with the mean value of 60 cem™. as estimated for the whole N66 region by Tsamisetal.(2003) using the above-mentioned ratio of the sulphur doublet. as well as a density of 80 cm? for N66 reported by Reidetal.(2006).," This agrees well with the mean value of 60 $^{-3}$, as estimated for the whole N66 region by \citet[][]{Tsamis03} using the above-mentioned ratio of the sulphur doublet, as well as a density of 80 $^{-3}$ for N66 reported by \citet[][]{Reid06}."
. Using the line ratio 5537/5517. which identifies denser zones. a density of 3700 em is found for N66 (Tsamisetal.2003).," Using the line ratio 5537/5517, which identifies denser zones, a density of 3700 $^{-3}$ is found for N66 \citep[][]{Tsamis03}."
 Furthermore. the total mass of the tonized gas. calculated from the with the previously noted Str6mmeren sphere assumption is presented in Col.," Furthermore, the total mass of the ionized gas, calculated from the with the previously noted Strömmgren sphere assumption is presented in Col."
 7., 7.
 The ionization is produced by Lyman continuum photon flux given in Col., The ionization is produced by Lyman continuum photon flux given in Col.
 8., 8.
 The most extincted part of the N66A region ts the dust lane (Sect., The most extincted part of the N66A region is the dust lane (Sect.
" 3.1). where the rratio reaches a value of 5.0 corresponding to A, = 1.7 mag."," 3.1), where the ratio reaches a value of 5.0 corresponding to $_{v}$ = 1.7 mag."
 Outside the lane. the ratio is smaller in particular around the main exciting stars.," Outside the lane, the ratio is smaller in particular around the main exciting stars."
" The average value of the Balmer decrement towards N66A is about 3.7 (A, = 0.8 mag).", The average value of the Balmer decrement towards N66A is about 3.7 $_{v}$ = 0.8 mag).
 We note that this ratio is higher than towards other areas of the rregion N66., We note that this ratio is higher than towards other areas of the region N66.
 As for the rratio. it fluctuates around 4.5 and reaches at the north-eastern border of N66À the value of 5.," As for the ratio, it fluctuates around 4.5 and reaches at the north-eastern border of N66A the value of 5."
 Spectral classification was derived for five massive. stars towards N66A. These stars were extracted from long-slit spectra., Spectral classification was derived for five massive stars towards N66A. These stars were extracted from long-slit spectra.
 Four spectra belong to grating #112 (N66A-1. MPG 455. MPG 595. MPG 655) and one spectrum (MPG 445) to erating #44.," Four spectra belong to grating 12 (N66A-1, MPG 455, MPG 595, MPG 655) and one spectrum (MPG 445) to grating 4."
 The identification of stars along the slits were based on monitor sketches drawn during the observations., The identification of stars along the slits were based on monitor sketches drawn during the observations.
 The spectral classification was performed using the criteria stated by Walborn&Fitzpatrick(1990)., The spectral classification was performed using the criteria stated by \citet[][]{Walborn90}.
. The rectified spectrograms. corrected for the nebulosity background. are displayed in Fig. 4..," The rectified spectrograms, corrected for the nebulosity background, are displayed in Fig. \ref{fig:sp}."
 The results are summarized in Table 3.. which also gives the corresponding astrometric and photometric information as well as classifications from previous studies when available.," The results are summarized in Table \ref{table:sp}, which also gives the corresponding astrometric and photometric information as well as classifications from previous studies when available."
 We note that the broader lines of 4445 is an instrumental effect (grating #44)., We note that the broader lines of 445 is an instrumental effect (grating 4).
 A particularly interesting star in this study. 1.6. N66À-1. was classified as VV. This spectral type agrees well with that based on the stellar Lyman continuum derived from the emission of the rregion.," A particularly interesting star in this study, i.e. N66A-1, was classified as V. This spectral type agrees well with that based on the stellar Lyman continuum derived from the emission of the region."
 However. the latter result is probably an underestimate because some of the ionizing photons were missed.," However, the latter result is probably an underestimate because some of the ionizing photons were missed."
 In particular. local dust absorbs a portion of the photons and. in addition. since the rregion is not fully ionization-bounded. some of the Lyman continuum photons escape into the interstellar medium.," In particular, local dust absorbs a portion of the photons and, in addition, since the region is not fully ionization-bounded, some of the Lyman continuum photons escape into the interstellar medium."
 This implies that the spectral classificatior based on spectroscopy is also an underestimate., This implies that the spectral classification based on spectroscopy is also an underestimate.
 This is probably because the spatial resolution of the spectra does not natch that of theHST ACS images. which show that the naim exciting stars are separated by 07.7.," This is probably because the spatial resolution of the spectra does not match that of the ACS images, which show that the main exciting stars are separated by .7."
 In other words. the spectrum of N66A-1 is contaminated by N66-2.," In other words, the spectrum of N66A-1 is contaminated by N66-2."
 If star #22 were of later type than star 311. the latter would be earlier than V. Νό6Α is clearly the most compact rregion of the N66 complex in the optical.," If star 2 were of later type than star 1, the latter would be earlier than V. N66A is clearly the most compact region of the N66 complex in the optical."
 [ts relative compactness. brightness. and location suggest that it is probably a relatively younger generation in the N66 complex.," Its relative compactness, brightness, and location suggest that it is probably a relatively younger generation in the N66 complex."
 Itshould belong to a distinct and rare class of regions in the Magellanic Clouds (MCs) called “Blobs”. or HEBs," Itshould belong to a distinct and rare class of regions in the Magellanic Clouds (MCs) called High-Excitation “Blobs”, or HEBs"
On the other hand. we can obtain an analytical estimate.' 71 of the turnover. frequeney for the hybrid electron. distribution. by considering a purely power-law electron distribution normalized to match the Maxwellian at u.,"On the other hand, we can obtain an analytical estimate, $\nu_{\rm t}^{\rm
pl}$, of the turnover frequency for the hybrid electron distribution by considering a purely power-law electron distribution normalized to match the Maxwellian at $\gamma_{\rm nth}$."
 For the svnchrotron absorption coellicient for. law electrons (Rvbicki Lightman 1979). which. with equation (2).. leads to where rods the classical electron racius and E is Eulers gamma function.," For the synchrotron absorption coefficient for power-law electrons (Rybicki Lightman 1979), which, with equation , leads to where $r_{\rm e}$ is the classical electron radius and $\Gamma$ is Euler's gamma function."
 Provided pzx3 (so that radiation from electrons with ~254 is negligible) and emission at £juth is dominated by non-thermal electrons (Le. 5;7541sam. Which condition can be checked a posteriori by comparing the source function of power-law electrons. to the Ravleigh-Jeans intensity). IAuthcdpU.," Provided $p \ga 3$ (so that radiation from electrons with $\gamma > \gamma_{\rm
f}$ is negligible) and emission at $\nnth$ is dominated by non-thermal electrons (i.e. $\gamma_{\rm f} > \gamma_{\rm t} > \gamma_{\rm nth}$, which condition can be checked a posteriori by comparing the source function of power-law electrons to the Rayleigh-Jeans intensity), $\nnth\simeq \nu_{\rm
t}^{\rm pl}$."
 On the other hand. +D/jh when emission from⋅ thermal electrons is. significant∢⊀⋅ at zrn and. in. general. The accuracy of this approximation is typically <10 »er cent (provided 74«5; and pz 3).," On the other hand, $\nu_{\rm t}^{\rm pl}<\nu_{\rm t}^{\rm th}$ when emission from thermal electrons is significant at $\nnth$ and, in general, The accuracy of this approximation is typically $\la 10$ per cent (provided $\gamma_{\rm t}<\gamma_{\rm f}$ and $p \ga 3$ )."
 See WZOO0 for a discussion. of⋅ approximations. of νι., See WZ00 for a discussion of approximations of $\ntth$.
" We⇁ can obtain: the dependences of entvs.""ytd on the plasma parameters by consideringpos the ratio. pPLyἐνth cusgiven w approximations and (7)].", We can obtain the dependences of $\nu_{\rm t}^{\rm nth}/\ntth$ on the plasma parameters by considering the ratio $\nu_{\rm t}^{\rm pl}/\ntth$ [given by approximations and ].
" First. at constant. 9. this ratio decreases with increasing O since then 4, increases with O see equation (9)."," First, at constant $\delta$, this ratio decreases with increasing $\Theta$ since then $\gamma_{\rm nth}$ increases with $\Theta$ [see equation ]."
 The dependence on 7p is stronger or £e than for e as a result of much steeper optically hin svnehrotron spectrum in the latter case., The dependence on $\tau_{\rm T}$ is stronger for $\nu_{\rm t}^{\rm pl}$ than for $\ntth$ as a result of much steeper optically thin synchrotron spectrum in the latter case.
" This leads to ertTu.TDchopUQUL Jie. the larger the optical depth. the arger the ratio: p""ATIfe."," This leads to $\nu_{\rm t}^{\rm pl}/\ntth \propto \tau_{\rm
T}^{2/(4+p)-0.05}$, i.e. the larger the optical depth, the larger the ratio $\nu_{\rm t}^{\rm pl}/\ntth$."
 In a hot accretion: How. we generally expect rp increasing and © decreasing with increasing accretion rate. so the ellect of a non-thermal tail (of à given orm) on the value of the turnover frequency can be expected o be the largest for the most luminous sources.," In a hot accretion flow, we generally expect $\tau_{\rm T}$ increasing and $\Theta$ decreasing with increasing accretion rate, so the effect of a non-thermal tail (of a given form) on the value of the turnover frequency can be expected to be the largest for the most luminous sources."
" Since (wqplfen""0XQ2ΑΕΙp)LOL Ίο non-thermal increase of the turnover frequency. grows with decreasing vo."," Since $\nu_{\rm t}^{\rm pl}/\ntth \propto \nu_{\rm c}^{(2+p)/(4+p)-0.91}$, the non-thermal increase of the turnover frequency grows with decreasing $\nu_{\rm c}$."
" This effect is due to the increase of 5, with decreasing vo -as x result of e.increasing. £47tl/£4. see equation. (6)]]. which. leads to a much larger ratio of the number of non-thermal to thermal electrons at 5, for a given 54."," This effect is due to the increase of $\gamma_{\rm t}$ with decreasing $\nu_{\rm c}$ [as a result of increasing $\ntth/\nu_{\rm c}$, see equation ], which leads to a much larger ratio of the number of non-thermal to thermal electrons at $\gamma_{\rm t}$ for a given $\gamma_{\rm
nth}$."
 Fhus. the elfect ofa non-thermal tail with a given form will be much stronger in AGNs than in BIIBs (due to much weaker magnetic field in the former objects).," Thus, the effect ofa non-thermal tail with a given form will be much stronger in AGNs than in BHBs (due to much weaker magnetic field in the former objects)."
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿∖, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿∖⋎, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿∖⋎⋯, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿∖⋎⋯↥, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
 ⋅↥↥⊔ ⋅: ↓⋅⋖⋅∐∶↓⋅⋜⊔↓∪⊳∖↓↥∪∖∖⊽⊳∖↿↓↥⋖⋅∖⇁⋜↧⇂⋯⊾⊳∖∪⇂∕∕↕⊔∆∕∕↥∆⋜↧⊳∖⋜↧⇂⊔⊔≼∼↿↓∪⊔∪⇂↿∖⋎⋯↥∆, \\ref{f:ratio} shows the values of $\nnth/\ntth$ as a function of
axisymmetric spacetime with a matter source (Bonazzolaetal.1993;Gourgoulhon 1999).,"axisymmetric spacetime with a matter source \citep{bonazzola1993,gourgoulhon1999}."
. The code has been tested extensively and compared with a few different numerical codes (Nozawaetal.1998)., The code has been tested extensively and compared with a few different numerical codes \citep{nozawa1998}.
". As theoretical calculations for dense matter at supranuclear densities are poorly constrained, the EOS in the high-density core of compact stars is not well understood (see, e.g., Weberetal.(2007);Haenselal.(2007) for reviews)."," As theoretical calculations for dense matter at supranuclear densities are poorly constrained, the EOS in the high-density core of compact stars is not well understood (see, e.g., \citet{weber2007,haensel2007} for reviews)."
" In this work, we employ eight realistic nuclear matter EOS to modelrotating neutron stars: model A 1971), model APR (Akmaletal. 1998),, (Pandharipandemodel AU (the AV14+UVII model in Wiringaetal.(1988) is joined to Negele&Vau- model BBB2 (Baldoetal.1997),, model FPS (Pandharipande(1973))),&Ravenhall1989;Lorenzetal. 1993),, model SLY4 (Douchin&Haensel2000),, model UU UVIA4--UVII model in Wiringaetal. is joined(the to Negele&Vautherin (1973))) and (1988)model WS (the UV14+TNI model in Wiringaetal.(1988) is joined to Lorenzetal. (1993)))."," In this work, we employ eight realistic nuclear matter EOS to modelrotating neutron stars: model A \citep{pandha1971}, model APR \citep{akmal1998}, , model AU (the AV14+UVII model in \citet{wiringa1988} is joined to \citet{negele1973}) ), model BBB2 \citep{baldo1997}, model FPS \citep{pandha1989,lorenz1993}, , model SLY4 \citep{douchin2000}, model UU (the UV14+UVII model in \citet{wiringa1988} is joined to \citet{negele1973}) ) and model WS (the UV14+TNI model in \citet{wiringa1988} is joined to \citet{lorenz1993}) )."
" For the quark star models, we use the simplest MIT bag model with non-interacting massless quarks (Chodosetal.1974)."," For the quark star models, we use the simplest MIT bag model with non-interacting massless quarks \citep{chodos1974}."
". Two different values, 60MeVfm? and 90MeVfm?, are chosen for the bag constant B."," Two different values, $60\ {\rm MeV\ fm}^{-3}$ and $90\ {\rm MeV\ fm}^{-3}$, are chosen for the bag constant $B$."
 These values correspond approximately to the range of B within which the hypothesis of strange matter is valid (Haenseletal. 2007)., These values correspond approximately to the range of $B$ within which the hypothesis of strange matter is valid \citep{haensel2007}.
". To illustrate the diversity of the EOS models used in this work, we plot the gravitational mass M against central energy density p. for non-rotating stars constructed with the chosen EOS models in Figure 1.."," To illustrate the diversity of the EOS models used in this work, we plot the gravitational mass $M$ against central energy density $\rho_{\rm c}$ for non-rotating stars constructed with the chosen EOS models in Figure \ref{fig:tov_curve}."
" The quark star models QMB60 and QMB90 in the figure correspond to the cases B=60MeVfm? and 90MeVfm?, respectively."," The quark star models QMB60 and QMB90 in the figure correspond to the cases $B=60\ {\rm MeV\ fm}^{-3}$ and $90\ {\rm MeV\ fm}^{-3}$, respectively."
 The maximum mass of compact stars depends quite sensitively on the EOS models and it ranges from about 1.5Mo to 2.2 , The maximum mass of compact stars depends quite sensitively on the EOS models and it ranges from about $1.5\ M_{\odot}$ to $2.2\ M_{\odot}$.
"In Figure 2,, we plot the Keplerian frequency fy against Mo.the gravitational mass M of rotating compact stars based on our chosen EOS models."," In Figure \ref{fig:fk_mass}, we plot the Keplerian frequency $f_{\rm K}$ against the gravitational mass $M$ of rotating compact stars based on our chosen EOS models."
" Similar to the maximum mass for non-rotating compact stars, Figure 2. shows clearly that fk depends strongly on the EOS model."," Similar to the maximum mass for non-rotating compact stars, Figure \ref{fig:fk_mass} shows clearly that $f_{\rm K}$ depends strongly on the EOS model."
 It is also sensitive to the mass of the star., It is also sensitive to the mass of the star.
 This is the well-known reason why searching for rapidly rotating compact stars can provide us constraints on the EOS of dense matter., This is the well-known reason why searching for rapidly rotating compact stars can provide us constraints on the EOS of dense matter.
 Now we turn to the main focus of this work: the dimensionless spin parameter j., Now we turn to the main focus of this work: the dimensionless spin parameter $j$.
" Having seen that fx depends strongly on the EOS and the mass of the star, it may be quite surprising to learn that the maximum value of the spin parameter jmax (as set by the Kepler limit) is quite universal for rotating neutron stars."," Having seen that $f_{\rm K}$ depends strongly on the EOS and the mass of the star, it may be quite surprising to learn that the maximum value of the spin parameter $j_{\rm max}$ (as set by the Kepler limit) is quite universal for rotating neutron stars."
" In Figure 3,, we plot jmax against the gravitational mass M for the selected nuclear matter EOS models."," In Figure \ref{fig:jmax_m2_ns}, we plot $j_{\rm max}$ against the gravitational mass $M$ for the selected nuclear matter EOS models."
" In the figure, each line corresponds to one particular EOS and each point on a line corresponds to a star model with a fixed M rotating at its Keplerian frequency."," In the figure, each line corresponds to one particular EOS and each point on a line corresponds to a star model with a fixed $M$ rotating at its Keplerian frequency."
 Note that each sequence in the figure is terminated at the stellar model with the same total particle number as the stable maximum-mass non-rotating configuration., Note that each sequence in the figure is terminated at the stellar model with the same total particle number as the stable maximum-mass non-rotating configuration.
" In the figure, we see that Jmax lies in a narrow range ~0.65—0.7 for the eight different nuclear matter EOS models."," In the figure, we see that $j_{\rm max}$ lies in a narrow range $\sim 0.65-0.7$ for the eight different nuclear matter EOS models."
" In particular, the values do not depend sensitively on the mass of the star."," In particular, the values do not depend sensitively on the mass of the star."
 This extends previous works (Cooketal.1994;Sal-gadoetal.1994) which focus on maximum-mass neutron star models.," This extends previous works \citep{cook1994,salgado1994}
 which focus on maximum-mass neutron star models."
" While the spin parameter of an astrophysical black hole is constrained by 7<1, we find that the spin parameter of a neutron star is boundedby jmax~0.7."," While the spin parameter of an astrophysical black hole is constrained by $j \le 1$, we find that the spin parameter of a neutron star is boundedby $j_{\rm max} \sim 0.7$."
 The upper bound jmax is quite universal for different EOS models and gravitational mass larger than ~1Mo., The upper bound $j_{\rm max}$ is quite universal for different EOS models and gravitational mass larger than $\sim 1\ M_{\odot}$.
" For lower mass neutron stars, M«1Mo, we find that j decreases with decreasing M."," For lower mass neutron stars, $M < 1\ M_\odot$, we find that $j$ decreases with decreasing $M$."
" However, we shall only focus on mass M>1Mg in this work as the observed masses of neutron stars are typically larger than 1Mo."," However, we shall only focus on mass $M > 1 \ M_\odot$ in this work as the observed masses of neutron stars are typically larger than $1\ M_\odot$."
 We refer the reader to Steineretal.(2010) for a recent review on the observed masses of neutron stars., We refer the reader to \citet{steiner2010} for a recent review on the observed masses of neutron stars.
 So far we have studied the maximum spin parameter jmax Of neutron stars rotating at their Keplerian frequencies fx., So far we have studied the maximum spin parameter $j_{\rm max}$ of neutron stars rotating at their Keplerian frequencies $f_{\rm K}$ .
" However, realistic neutron stars in general rotate slower with frequencies f« fx."," However, realistic neutron stars in general rotate slower with frequencies $f < f_{\rm K}$ ."
 Is thespin parameter jstill insensitive to the EOS and mass of the star?, Is thespin parameter $j$still insensitive to the EOS and mass of the star?
" In Figure 4,, we plot the spin parameter j "," In Figure \ref{fig:FPS_j_f}, , we plot the spin parameter $j$ "
The GIS and SIS spectra thus obtained are illustrated in Figure 4..,The GIS and SIS spectra thus obtained are illustrated in Figure \ref{fig:ascaspec}.
 Following bwasawa Comastri (1998). we fit the spectrum of NGC 6240 with à two-component model that includes a thermal component and an AGN component.," Following Iwasawa Comastri (1998), we fit the spectrum of NGC 6240 with a two-component model that includes a thermal component and an AGN component."
 The thermal component consists of emission. [rom two optically-thin thermal plasmas of different temperatures. where the higher temperature. component has an excess absorption.," The thermal component consists of emission from two optically-thin thermal plasmas of different temperatures, where the higher temperature component has an excess absorption."
 The AGN component consists of an absorbed power-law continuum and line emission., The AGN component consists of an absorbed power-law continuum and line emission.
 Phe mocel can be wrltten as: where the parameters in parentheses are kept free., The model can be written as; where the parameters in parentheses are kept free.
 For the thermal plasma emission. code. we employ. the Alewe-Ixaastra model (Alewe. Gronenschild. van den Oord 1985: Mewe. Lemen. van den Oord. 1986: Ixaastra. 1092) mocüfied by οσα. Osterheld. Goldstein (1995). which is implemented. in ASPEC as the MEIAL model.," For the thermal plasma emission code, we employ the Mewe-Kaastra model (Mewe, Gronenschild, van den Oord 1985; Mewe, Lemen, van den Oord 1986; Kaastra 1992) modified by Liedahl, Osterheld, Goldstein (1995), which is implemented in XSPEC as the MEKAL model."
 The metallicities. Z. o£the two thermal components are assumed to be the same.," The metallicities, $Z$, of the two thermal components are assumed to be the same."
 The abundance ratios among clferent elements are fixed to be the fiducial solar values given bv Anders Crevesse (1989)., The abundance ratios among different elements are fixed to be the fiducial solar values given by Anders Grevesse (1989).
 With this model. we fit the GIS ancl SIS. spectra simultaneously.," With this model, we fit the GIS and SIS spectra simultaneously."
 In this simultaneous fit. the CIS data below 2.0 keV are excluded because of a slight uncertainty in the energy scale of the CS arising from the complex xenon. edge structure in this range.," In this simultaneous fit, the GIS data below 2.0 keV are excluded because of a slight uncertainty in the energy scale of the GIS arising from the complex xenon M-edge structure in this range."
 A good Lit is obtained as shown in Figure 4 and Table 2. and the results are essentially the same as obtained by lwasawa Comastri (1998) using the SIS data only.," A good fit is obtained as shown in Figure \ref{fig:ascaspec} and Table 2, and the results are essentially the same as obtained by Iwasawa Comastri (1998) using the SIS data only."
 As noted by Iwasawa Comastri (1998). the striking spectral features are a very [lat continuum. above  keV with a photon index of 0 and a strong iron Ix-line with an equivalent width of 1.2 keV. Although the contamination sources near NGC 6240 are not identified. we only need the energy spectrum for the purpose of the present work.," As noted by Iwasawa Comastri (1998), the striking spectral features are a very flat continuum above $\sim$ 4 keV with a photon index of $\sim 0$ and a strong iron K-line with an equivalent width of 1.2 keV. Although the contamination sources near NGC 6240 are not identified, we only need the energy spectrum for the purpose of the present work."
 We construct the spectrum from the whole field of view of the GIS excluding the 3-aremin radius circle centered on the X-ray peak of NGC 6240. from which the estimated background (NXD. | CNB) is subtracted.," We construct the spectrum from the whole field of view of the GIS excluding the 3-arcmin radius circle centered on the X-ray peak of NGC 6240, from which the estimated background (NXB + CXB) is subtracted."
 In aciedition. because of an extended outskirts of the XI point-spread function. the spectrum still contains a significant contribution from NGC 6240.," In addition, because of an extended outskirts of the XRT point-spread function, the spectrum still contains a significant contribution from NGC 6240."
 Using a ray-tracing simulation software. the contribution from NGC 6240 outside the 3-arcmin radius circle is estimated and subtracted.," Using a ray-tracing simulation software, the contribution from NGC 6240 outside the 3-arcmin radius circle is estimated and subtracted."
 Figure 5 shows the GLS spectrum thus obtained., Figure \ref{fig:contamispec} shows the GIS spectrum thus obtained.
 The resultant CS spectrum contains the total photons from the contamination sources within the CIS field of view and the Galactic soft ΠΠ emission., The resultant GIS spectrum contains the total photons from the contamination sources within the GIS field of view and the Galactic soft diffuse emission.
 As shown in Figureὃν 5.. thus obtained. GIS spectrum is fittecl satisfactorily QCv= 184/25) with the sum ofa thin- model (ALEKAL model) and à power-law model.," As shown in Figure \ref{fig:contamispec}, thus obtained GIS spectrum is fitted satisfactorily $\chi^2/\nu=18.4/25$ ) with the sum of a thin-thermal model (MEKAL model) and a power-law model."
 For, For
"and Lastly, using both of these estimates, the (numerous) parameters governing the jet luminosity may be constrained: where [ολο is the peak luminosity of the jet in units of 10*° erg s.","and Lastly, using both of these estimates, the (numerous) parameters governing the jet luminosity may be constrained: where $L_{jet,46}$ is the peak luminosity of the jet in units of $10^{46}$ erg $^{-1}$."
 The recently-discovered flare source J2058.4--0516 (Cenkoetal.2011) may be an example of exactly the sort of tidal disruption event to which this formalism is applicable., The recently-discovered flare source J2058.4+0516 \citep{cenko11} may be an example of exactly the sort of tidal disruption event to which this formalism is applicable.
 Its lightcurve is shown in Figure 3.., Its lightcurve is shown in Figure \ref{fig:lc}.
" For ~10 d, its flux stayed nearly constant; for the next three months, the flux declined roughly as a power-law in time."," For $\simeq 10$ d, its flux stayed nearly constant; for the next three months, the flux declined roughly as a power-law in time."
" Around t~100 d, the decline became shallower; unfortunately, the flux at that point was already near Swift’s"," Around $t \sim 100$ d, the decline became shallower; unfortunately, the flux at that point was already near Swift's"
"The object commonly known as Barnard's Loop was discovered photographically more than a century ago, originally by W. H. Pickering in 1889 (Sheehan1995) and then by E. E. Barnard(1894).","The object commonly known as Barnard's Loop was discovered photographically more than a century ago, originally by W. H. Pickering in 1889 \citep{she95} and then by E. E. \citet{bar94}."
". Although a popular object for imaging by amateur astronomers, its low surface brightness has delayed the definitive observations necessary for developing a good model of its characteristics."," Although a popular object for imaging by amateur astronomers, its low surface brightness has delayed the definitive observations necessary for developing a good model of its characteristics."
 A representative image is shown in Figure 1., A representative image is shown in Figure 1.
 It is an arc of about facing east and at a distance of 440 pc (O'Dell&Henney2008) is 110 pc in length., It is an arc of about facing east and at a distance of 440 pc \citep{oh08} is 110 pc in length.
 The northern boundary is at a Galactic Latitude of aand lines of galactic longitude are at a position angle of about62°., The northern boundary is at a Galactic Latitude of and lines of galactic longitude are at a position angle of about.
. These facts mean that its two sides (north and south) extend over a distance of about 60 pc and the conditions of the local ambient interstellar medium will have changed significantly., These facts mean that its two sides (north and south) extend over a distance of about 60 pc and the conditions of the local ambient interstellar medium will have changed significantly.
" The north edge is about twice as bright as the south, indicating a higher density there."," The north edge is about twice as bright as the south, indicating a higher density there."
 The structure of Barnard's Loop was the subject of a study of the brightness distribution of ultraviolet light measured in a photograph obtained during the Gemini 11 manned space-flight., The structure of Barnard's Loop was the subject of a study of the brightness distribution of ultraviolet light measured in a photograph obtained during the Gemini 11 manned space-flight.
 In this investigation O'Delletal.(1967) derived the distribution of interstellar dust under the assumption that the ultraviolet continuum is due to scattered light originating in the OB stars in the Orion constellation Belt and Sword regions., In this investigation \citet{oyh67} derived the distribution of interstellar dust under the assumption that the ultraviolet continuum is due to scattered light originating in the OB stars in the Orion constellation Belt and Sword regions.
 It was, It was
is responsible for its ireeular shape).,is responsible for its irregular shape).
 We sec that the N-vay spectra implied bv these data is very simular to that of typical Sevtert 1s. with [TyL8 2.0.," We see that the X-ray spectrum implied by these data is very similar to that of typical Seyfert 1s, with $\Gamma_{\rm X}\sim 1.8$ –2.0."
 The obtained range of teniperature is relatively low. AL~ 5080 keV. iu aerecmicut with the results of Zdaziarski et ((1996) and J97.," The obtained range of temperature is relatively low, $kT\sim 50$ –80 keV, in agreement with the results of Zdziarski et (1996) and J97."
 The weighted average of the 50.200 keV iudices of individual objects (Table 1 aud Figure 1)) and he iudices in the power-law fits to tho stu spectra (Table 2) for Sevtert Ls. 2.5040.09 and 2.56+0.11. respectively. are significantly: softer than those or Sevfert 2s. 2.05+0.09 aud 2.21+0.12.," The weighted average of the 50–200 keV indices of individual objects (Table 1 and Figure \ref{indices}) ) and the indices in the power-law fits to the sum spectra (Table 2) for Seyfert 1s, $2.50\pm 0.09$ and $2.56\pm 0.14$, respectively, are significantly softer than those for Seyfert 2s, $2.05\pm 
0.09$ and $2.21\pm 0.12$."
 The xobabilitv that the 2 samples are drawn from the sale distribution is oulv (82))., The probability that the 2 samples are drawn from the same distribution is only \ref{s:data}) ).
 This difference is confirmed by Figures 3. LL. where we see that he averageOo spectrum of Sevfert 2s is noticeabel varder than that of Sevfert Ls.," This difference is confirmed by Figures \ref{spectra}, \ref{contours}, where we see that the average spectrum of Seyfert 2s is noticeably harder than that of Seyfert 1s."
 We lave investigated whether this difference can be explained by the viewing angle different )etween Sevtert ls and 2s., We have investigated whether this difference can be explained by the viewing angle different between Seyfert 1s and 2s.
 One relevaut effect is he streneth of Compton reflection decreasing with he increasing viewiug angele., One relevant effect is the strength of Compton reflection decreasing with the increasing viewing angle.
 Since the spectrum roni Compton reflection typically peaks around 30 seV. followed by a steep decline at higher energies. he larger R the softer the spectrum iu the OSSE range.," Since the spectrum from Compton reflection typically peaks around 30 keV followed by a steep decline at higher energies, the larger $R$ the softer the spectrum in the OSSE range."
 However. this effect is already inchided iu our Comptonizatiou/reficction model. aud Figure 1l. shows that it is not sufücient to account for the difference between the g puariuueters of the average spectra.," However, this effect is already included in our Comptonization/reflection model, and Figure \ref{contours} shows that it is not sufficient to account for the difference between the $y$ parameters of the average spectra."
 We note. however. that the average inclination of Sevfert 28 remains unknown.," We note, however, that the average inclination of Seyfert 2s remains unknown."
 We have found we can fit the two spectra with the same Coniptouization/roflection model (vith R= 0.75) if cos;=O!8° for Seyfert 2s within confidence., We have found we can fit the two spectra with the same Comptonization/reflection model (with $R=0.75$ ) if $\cos i=0^{+0.3}$ for Seyfert 2s within confidence.
 Aun additional subtle effect appears when the Comptonizine inediun las ao slab geometry., An additional subtle effect appears when the Comptonizing medium has a slab geometry.
 Naiuely. photous euütted at a large viewing angele (with respect to the slab normal) undergo a arecr nuniber of scatterings than those enuütted at a sinall viewing angle. due to the escape xobabilitv at a eiven depth from the surface. 7. ina eiven direction being exptz//cos7}.," Namely, photons emitted at a large viewing angle (with respect to the slab normal) undergo a larger number of scatterings than those emitted at a small viewing angle, due to the escape probability at a given depth from the surface, $\tau'$, in a given direction being $\exp(-\tau'/\cos i)$."
 We rave thus fitted our data with a model in which he Comptonizing medium forms a slab with a ralfthickuess 7 aud the seed pliotous are emitted Ww point sources m the midplane., We have thus fitted our data with a model in which the Comptonizing medium forms a slab with a half-thickness $\tau$ and the seed photons are emitted by point sources in the midplane.
" We have first fitted with this model the Ginye-OSSE spectruu of (96, and then fixed the obtained Π iu the fits to our average spectra."," We have first fitted with this model the -OSSE spectrum of G96, and then fixed the obtained $R$ in the fits to our average spectra."
 The resulting best-fit yaralneters are AE—63 keV. 07 keV. g=0.21. V.18. for Sevtert Ls aud 2s. respectively.," The resulting best-fit parameters are $kT=63$ keV, 67 keV, $y=0.34$, 0.48, for Seyfert 1s and 2s, respectively."
 The error contours are shown iu Fieure 5.., The error contours are shown in Figure \ref{slab}.
 We see that the wo paralcter ranges are still incompatible with each other., We see that the two parameter ranges are still incompatible with each other.
 This meaus that the spectral shape does not change sufficiently due to this effect o account for the observed spectral difference (sce Figure 3)), This means that the spectral shape does not change sufficiently due to this effect to account for the observed spectral difference (see Figure \ref{spectra}) ).
 For example. in the case of he best-fit Seytert-l model spectrum. (ποπιο voth Comptonization aud reflection). the ratio of monochromatic fluxes at 150 keV and 50 τον inereases oulv bv when the melinatiou changes from cos/=0.587 to cos/=0.L.," For example, in the case of the best-fit Seyfert-1 model spectrum (including both Comptonization and reflection), the ratio of monochromatic fluxes at 150 keV and 50 keV increases only by when the inclination changes from $\cos i=0.87$ to $\cos i=0.4$."
 The above results also show that the obtained values of AL and Ex only weakly depeud on ecolctry.," The above results also show that the obtained values of $kT$ and $\Gamma_{\rm
X}$ only weakly depend on geometry."
 For Sevfert 2s. our results imply Ex between Ll and 1.9 aud ET between 50 aud 17 keV. We then cousider possible effects of absorption/scatteriug in a torus simroundius the ceutral source beige inportaut in Sevtert 2s.," For Seyfert 2s, our results imply $\Gamma_{\rm X}$ between 1.4 and 1.9 and $kT$ between 50 and 170 keV. We then consider possible effects of absorption/scattering in a torus surrounding the central source being important in Seyfert 2s."
 This cau be an important effect in the OSSE range oulv when the torus is Thonsou-thick. Ng>15«10?! cm?.," This can be an important effect in the OSSE range only when the torus is Thomson-thick, $N_{\rm H} \ga 1.5\times 10^{24}$ $^{-2}$."
 Then the hardening occurs (apart from a weak effect of photo-clectric absorption) because the torus becomes lore transparent to scattering with the Increasing photon energy due to the Nleiu-Nishina effects., Then the hardening occurs (apart from a weak effect of photo-electric absorption) because the torus becomes more transparent to scattering with the increasing photon energy due to the Klein-Nishina effects.
 We have applied a muiuerical model of torus absorption/scatteriug of Irolik. Madau. Zvveki (0991) to our average spectu of Sevfer 2s assunniug the intrinsic spectrum from thenua Conmptouization of Sevfert 1s.," We have applied a numerical model of torus absorption/scattering of Krolik, Madau, Żyycki (1994) to our average spectrum of Seyfert 2s assuming the intrinsic spectrum from thermal Comptonization of Seyfert 1s."
 We have obtaine a good fitJ (AZD=“yeaye 22/26) at Ny-=23.«E105D cni* with an intrinsic TC spectrum icdeutica in the cases of Sevfert Is aud 2s with a torus with an opening angle of 60°., We have obtained a good fit $\chi_\nu^2=22/26$ ) at $N_{\rm H}=3\times 10^{24}$ $^{-2}$ with an intrinsic TC spectrum identical in the cases of Seyfert 1s and 2s with a torus with an opening angle of $60\degr$.
 IIowever. since that numerical model does not allow for includius Compton reflection from a disk. we camnot obtain mcanineful coustraiuts on its parameter space.," However, since that numerical model does not allow for including Compton reflection from a disk, we cannot obtain meaningful constraints on its parameter space."
 On the other haud. the ouly Seyfert 2 in our suuple with a Thomsou-thick absorber is NGC 1915. (Bisaliti. Maiolino Salvati 1999). with Nycdos1025 7o (see also Madejski et 22000). which uucleus is also viewed edge-on (Crecnhill. Moran Uerrnsteim 1997).," On the other hand, the only Seyfert 2 in our sample with a Thomson-thick absorber is NGC 4945 (Risaliti, Maiolino Salvati 1999), with $N_{\rm
H}\simeq 4\times 10^{24}$ $^{-2}$ (see also Madejski et 2000), which nucleus is also viewed edge-on (Greenhill, Moran Herrnstein 1997)."
 Also. some effect of absorption ou the OSSE baud is possible in Msn 3. with Nyyc1.3s107! cm7 (Cappi et," Also, some effect of absorption on the OSSE band is possible in Mkn 3, with $N_{\rm H}\simeq 1.3\times 10^{24}$ $^{-2}$ (Cappi et"
where € is the distance from the centre of the lens in the projected lens plane anc where a is the one-dimensional velocity dispersion. of ‘particles’ within the gravitational potential of the mass distribution. such as stars.,"where $\xi$ is the distance from the centre of the lens in the projected lens plane and where $\sigma_v$ is the one-dimensional velocity dispersion of `particles' within the gravitational potential of the mass distribution, such as stars."
" The dimensionless surface mass density or convergence is defined asa=MYM where X, (or the critical density) is clefinect as where D, and 2; are the angular cliameter distance to the source ancl lens. respectively. and 2), is the angular diameter distance between lens anc source."," The dimensionless surface mass density or convergence is defined as $\kappa = \Sigma/\Sigma_c$, where $\Sigma_c$ (or the critical density) is defined as where $D_s$ and $D_l$ are the angular diameter distance to the source and lens, respectively, and $D_{ls}$ is the angular diameter distance between lens and source."
 Thus for the case of the simple isothermal sphere we have where 6=€/)) is the angular clistance from lens centre in the sky plane and where @e is the Einstein dellection angle. defined as The llexion. J£. caused by the SIS at an angular vector displacement. 9. [rom the lens centre on the sky. plane is thus simply where 6 is the position angle around. the lens. and in this case also gives the direction of the flexion.," Thus for the case of the simple isothermal sphere we have where $\theta=\xi/D_l$ is the angular distance from lens centre in the sky plane and where $\theta_E$ is the Einstein deflection angle, defined as The flexion, $\flex$ , caused by the SIS at an angular vector displacement, $\thetab$, from the lens centre on the sky plane is thus simply where $\phi$ is the position angle around the lens, and in this case also gives the direction of the flexion."
 The first Ilexion for this profile is therefore circularly symmetric and (expressed as a vector) directed racially inwards towards the lens centre. as would be expected.," The first flexion for this profile is therefore circularly symmetric and (expressed as a vector) directed radially inwards towards the lens centre, as would be expected."
 Similarly. the second Iexion is: This has a larger maximum amplitude than the. first Ilexion for this lens profile. fades off with the same power law index away from the lens. and oscillates around the lens as a spin-3 quantity rather than a spin-1 quantity.," Similarly, the second flexion is: This has a larger maximum amplitude than the first flexion for this lens profile, fades off with the same power law index away from the lens, and oscillates around the lens as a spin-3 quantity rather than a spin-1 quantity."
 l]laving considered. the specific case of an isothermal sphere. we can continue more generally with power law representations of the shear around a lens: where cl is a constant. n=| corresponds to an isothermal sphere. n=2 corresponds to a point mass. and so on.," Having considered the specific case of an isothermal sphere, we can continue more generally with power law representations of the shear around a lens: where $A$ is a constant, $n=1$ corresponds to an isothermal sphere, $n=2$ corresponds to a point mass, and so on."
 In. particular. one can ask whether one can better describe the arced nature of lensed objects by the [lexion we have defined. or the shear derivatives themselves.," In particular, one can ask whether one can better describe the arced nature of lensed objects by the flexion we have defined, or the shear derivatives themselves."
 In order to answer this question. for simplicity we rotate the system such that the source lies along the |. axis from the lens.," In order to answer this question, for simplicity we rotate the system such that the source lies along the $+x$ axis from the lens."
" We then consider what the second-order lensing amplitudes would be in a 7derivative-space."" composed. of the (wo non-zero shear derivatives: In 7""lexion-space."" where the components are the first and second [exions. the second-order lensing amplitudes are: We wish to find out which is the most compact basis space."," We then consider what the second-order lensing amplitudes would be in a “derivative-space,” composed of the two non-zero shear derivatives: In “flexion-space,” where the components are the first and second flexions, the second-order lensing amplitudes are: We wish to find out which is the most compact basis space."
 For any given distribution. this will be the one for =vhich only one eigenstate is non-zero.," For any given distribution, this will be the one for which only one eigenstate is non-zero."
 ligure 2. shows the amplitudes in cach of these two μα»icees as a function of the shear power law index., Figure \ref{fg:eigenspace} shows the amplitudes in each of these two spaces as a function of the shear power law index.
 We see that. for point sources. Uexion space is the most. compact approach: the signal is a pure second flexion state.," We see that, for point sources, flexion space is the most compact approach; the signal is a pure second flexion state."
 For a galaxy profile with n21. both spaces are almost equally licient in describing the second order lensing.," For a galaxy profile with $n \simeq
1$, both spaces are almost equally efficient in describing the second order lensing."
 Adcitionally. »Xh the flexion and. derivative notations can be shown to »oduce 4 statistically independent terms. which. taken over an ensemble of images will all have mean zero.," Additionally, both the flexion and derivative notations can be shown to produce 4 statistically independent terms, which, taken over an ensemble of images will all have mean zero."
 Moreover. a representation. the standard. deviations of the two erms due to both intrinsic variation and photon noise wil »e identical.," Moreover, a representation, the standard deviations of the two terms due to both intrinsic variation and photon noise will be identical."
 We conclude that flexion. is an cllicient means of describing third order lensing., We conclude that flexion is an efficient means of describing third order lensing.
 For point masses it is optimal: or SIS galaxies it is as σους as considering shear derivatives: and in addition the division between local and. non-loca components which it exclusively allores is very. valuable., For point masses it is optimal; for SIS galaxies it is as good as considering shear derivatives; and in addition the division between local and non-local components which it exclusively affords is very valuable.
 1 also describes correctly the spin properties of the lensing., It also describes correctly the spin properties of the lensing.
 The SIS mass distribution can be modified so as to remove one feature whichmay not be a good. physical description, The SIS mass distribution can be modified so as to remove one feature whichmay not be a good physical description
ormation recipes.,formation recipes.
 lIlowever. for the purposes of the kinematic modelling. the cold gas need not be rotationally supported by disks we simply mean flattened: systems.," However, for the purposes of the kinematic modelling, the cold gas need not be rotationally supported — by disks we simply mean flattened systems."
 We emphasize that this description is a crude simplification of the distribution and dynamics of cold gas in real galaxies. especially in interacting galaxies. which show rather complex behavior (?)..," We emphasize that this description is a crude simplification of the distribution and dynamics of cold gas in real galaxies, especially in interacting galaxies, which show rather complex behavior \citep{hibb:00}."
 In the spirit of semi-analvtie models. it is our hope that such a simplified description can still capture the eencral behavior of the cold gas when averaged over a large number of svstenis.," In the spirit of semi-analytic models, it is our hope that such a simplified description can still capture the general behavior of the cold gas when averaged over a large number of systems."
 The high-ionization state absorption is assumed. to arise from the hot gas in halos., The high-ionization state absorption is assumed to arise from the hot gas in halos.
 In the standard galaxy. formation model. all the hot gas is assumed to be at a single temperature and only a single gas phase is considered.," In the standard galaxy formation model, all the hot gas is assumed to be at a single temperature and only a single gas phase is considered."
 These simplifications may be reasonable for calculating the cooling rate of the eas. but for describing kinematics we will need to consider the substructure of the hot gas. and we will assume Chat the majority of it is in a phase suitable. for CIV absorption. in general at temperatures less than the virial temperature of the halo.," These simplifications may be reasonable for calculating the cooling rate of the gas, but for describing kinematics we will need to consider the substructure of the hot gas, and we will assume that the majority of it is in a phase suitable for CIV absorption, in general at temperatures less than the virial temperature of the halo."
 We will continue to refer to this as hot gas because it is dvnamically hot even if parts of it have cooled., We will continue to refer to this as hot gas because it is dynamically hot even if parts of it have cooled.
 To compare to the observations we must convert. the hot gas mass toa CIV column density. ον.," To compare to the observations we must convert the hot gas mass to a CIV column density, $N_{CIV}$."
" Phe column density of CIV along a line of sight is related to the column density of hydrogen by Na. where Zi, is the metallicity of the gas and feay is the fraction of the gas in a state that produces CIV absorption."," The column density of CIV along a line of sight is related to the column density of hydrogen by , where $Z_{hg}$ is the metallicity of the gas and $\fciv$ is the fraction of the gas in a state that produces CIV absorption."
" Vhe column density in LIL is related to the total amount of hvdrogen by Ny,=(1wd)Ny. where ur is the ionization fraction of the gas."," The column density in HI is related to the total amount of hydrogen by $N_{HI}=(1-x)N_{H}$, where $x$ is the ionization fraction of the gas."
 The quantities c.fei and Zpig can vary locally in the gas. but we will assume that suitably averaged quantities can be defined to give the elobal relationships we use here.," The quantities $x,\fciv$ and $Z_{hg}$ can vary locally in the gas, but we will assume that suitably averaged quantities can be defined to give the global relationships we use here."
 We emphasize that only the metal lines in the systems are measured. which creates an inherit degeneracy between the total amount of gas. its metallicity ancl its ionization state.," We emphasize that only the metal lines in the systems are measured, which creates an inherit degeneracy between the total amount of gas, its metallicity and its ionization state."
 For simplicity we will treat the ionization state and metallicity to be uniform within cach halo and. therefore Neay is directly proportional to Nig., For simplicity we will treat the ionization state and metallicity to be uniform within each halo and therefore $N_{CIV}$ is directly proportional to $N_{H}$.
" The reader should bear in mind however that an increase in Wy can equally well be thought of as an increase in fegy or Z,,.", The reader should bear in mind however that an increase in $N_{H}$ can equally well be thought of as an increase in $\fciv$ or $Z_{hg}$ .
 We well return to this issue in 6.., We well return to this issue in \ref{sec:conc}.
 For simplicity we take fegy=1. noting that a smaller value of fegy can be directly offset by an increase in the hot gas metallicitv.," For simplicity we take $\fciv=1$, noting that a smaller value of $\fciv$ can be directly offset by an increase in the hot gas metallicity."
" We fix Zyi4 to be 1.5 dex Lor all hot eas in all halos. which produces a good match to the CIV column density distribution in DLA svstenis as measured in ον,"," We fix $Z_{hg}$ to be $-1.5$ dex for all hot gas in all halos, which produces a good match to the CIV column density distribution in DLA systems as measured in \citet{wp:00a}."
 The mean metallicity of DLA systems at z292 is xs 1.5to Ldex (?).. which is consistent with our assumed Zig.," The mean metallicity of DLA systems at $z > 2$ is $\approx -1.5$ to $-1$ dex \citep{pw:02}, which is consistent with our assumed $Z_{hg}$."
 Mowever. because of the degeneracy mentioned. above. this value is directly. proportional to our choice of feuy," However, because of the degeneracy mentioned above, this value is directly proportional to our choice of $\fciv$."
 ‘Thus this can be thought of as à minimum metallicity since foav<1., Thus this can be thought of as a minimum metallicity since $\fciv \le 1$.
" A realistic estimate of fe; might be z0.35 which would then require Z,4%189 dex."," A realistic estimate of $\fciv$ might be $\approx 0.35$ which would then require $Z_{hg}
\approx -1.0$ dex."
 Alternatively. the hot gas metallicity could be lower if there is more hot gas in halos than calculated by our simple recipes.," Alternatively, the hot gas metallicity could be lower if there is more hot gas in halos than calculated by our simple recipes."
 The ealaxy formation code calculates the metallicity of the hot gas as part of the chemical enrichment: mocel described. brielly above., The galaxy formation code calculates the metallicity of the hot gas as part of the chemical enrichment model described briefly above.
 We find that if we use. the metallicities calculated. in this way (which typically show a large scatter at fixed halo mass). we obtain too large a spread in the CIV. column density. distribution.," We find that if we use the metallicities calculated in this way (which typically show a large scatter at fixed halo mass), we obtain too large a spread in the CIV column density distribution."
 This is not too disturbing. as the modelling of the reheating and ejection of metals by supernova feedback. is extremely. uncertain.," This is not too disturbing, as the modelling of the reheating and ejection of metals by supernova feedback is extremely uncertain."
 The galaxy formation model was calibrated το produce the correct mean observed stellar. cold. gas. and hot. gas metallicities in 2=0 galaxies and clusters. but this is no guarantee that it produces the correct results for these quantities at 2=3. or the correct. dispersion in the niass-metallicity relation.," The galaxy formation model was calibrated to produce the correct mean observed stellar, cold gas, and hot gas metallicities in $z=0$ galaxies and clusters, but this is no guarantee that it produces the correct results for these quantities at $z=3$, or the correct dispersion in the mass-metallicity relation."
 Llowever. it is also possible that correlations between fou. Zig und the gas density conspire to reduce the spread in observed CIV. column densities.," However, it is also possible that correlations between $\fciv$, $Z_{hg}$ and the gas density conspire to reduce the spread in observed CIV column densities."
 An example of this kind of conspiracy would be a starburst which produces high metallicity in the hot gas but also reduces the hot gas column density or changes the ionization state of the gas such that CIV is no longer thepreferred. ionization state of carbon., An example of this kind of conspiracy would be a starburst which produces high metallicity in the hot gas but also reduces the hot gas column density or changes the ionization state of the gas such that CIV is no longer thepreferred ionization state of carbon.
 Unfortunately. the large number of uncertaintiesmean that [ew constraints can be placed on any aspect of the modelling that &oes into determining CIV column densities.," Unfortunately, the large number of uncertaintiesmean that few constraints can be placed on any aspect of the modelling that goes into determining CIV column densities."
 We discuss, We discuss
foot strength of >300 CG. have the other foot point iu regions with «300 €. Fieure Ub shows the photospheric vertical field |B.| statistics for pixels hosting a leading (stronger) foot poiut (solid histoeram-stvle line) aud for pixel hosting a weak foot point (dashed line}.,"foot strength of $>300$ G, have the other foot point in regions with $<300$ G. Figure \ref{fig4}b b shows the photospheric vertical field $|B_z|$ statistics for pixels hosting a leading (stronger) foot point (solid histogram-style line) and for pixel hosting a weak foot point (dashed line)."
 Obviously the vertical field of both foot points for long loops (veaching at least into the chromosphere) is very different., Obviously the vertical field of both foot points for long loops (reaching at least into the chromosphere) is very different.
 Partly. this las to do with the differeit field streneth distribution in the two magnetic polarities for this particular maguctogram.," Partly, this has to do with the different field strength distribution in the two magnetic polarities for this particular magnetogram."
 The anit vertical field iu the negative field region is 1311 C. compared to 707 € iu positive polarity PC@IOUS5," The maximum vertical field in the negative field region is $1311$ G, compared to $707$ G in positive polarity regions."
5 Tuvestigations of further quiet Sun regions with high resolution are necessary to see if such differeut field strengths are common or au exception (ForquietSuninvestigationsofmae-ucticfluxsec.e.g.Waneetal. 1995).," Investigations of further quiet Sun regions with high resolution are necessary to see if such different field strengths are common or an exception \citep[For quiet Sun investigations of magnetic flux see, e.g.][]{wang:etal95}."
. We investigated the simallscale structure of nagnetic fields im the solu atmosphere with emphasis on the magnetic connectivity between i6 photosphere and chromosphere., We investigated the small-scale structure of magnetic fields in the solar atmosphere with emphasis on the magnetic connectivity between the photosphere and chromosphere.
 We fore iat chromospheric areas σουαπο about 91:4 of 160 Inagnetic euerev are topologically councecte o a phDotosphneric foot poimt with a streneth Habove 300 αν Le. with an energv deusitv higher wan the average equipartition with that of convective flows.," We found that chromospheric areas containing about $91 \%$ of the magnetic energy are topologically connected to a photospheric foot point with a strength above $300$ G, i.e. with an energy density higher than the average equipartition with that of convective flows."
 For the majority of the loops ie magnetic field strength iu both foot points oc.ffe significauth. with the second foot poii raving a field streneth below equipartition. ic. retwork clemeuts connect magnetically mainly to internetwork (IN) features. in general agreenmienu with uunernceal experinueuts of Schrier&Title(2003):Jeudersie&Peter (2006).," For the majority of the loops the magnetic field strength in both foot points differs significantly, with the second foot point having a field strength below equipartition, i.e. network elements connect magnetically mainly to internetwork (IN) features, in general agreement with numerical experiments of \cite{schrijver:etal03a,jendersie:etal06}."
. Au interestingo estion is fo which extent the fudiug tha network clements are connected mainly with IN features wight infiacuce models of quict-Sun loop heating (e.g.Tanstecn1993:Chaeetal.2002:Milleretal.2003. 2001).," An interesting question is to which extent the finding that network elements are connected mainly with IN features might influence models of quiet-Sun loop heating \citep[e.g.,][]{hansteen93,chae:etal02,mueller:etal03,mueller:etal04}."
. Such loop heating models usually assiunie a constant cross section. al approximation which is not fulfilled for field lines connecting strong and weak field regions in the photosphere.," Such loop heating models usually assume a constant cross section, an approximation which is not fulfilled for field lines connecting strong and weak field regions in the photosphere."
 It isalso worth noting that IN fields are more dvnanue aud short-lived than network fields (INfieldaveragelifetimesareabout10miuaccordingtodeWijunetal.2008).," It isalso worth noting that IN fields are more dynamic and short-lived than network fields \citep[IN field average lifetimes are about 10 min according
to][]{dewijn:etal08}."
. Since nost chromospheric aud coronal loops. at least iu the observed region. are connected at one foot point with IN fields we expect the chromospheric aud coronal field to be more dynamic than sugeested by uetwork fields alone.," Since most chromospheric and coronal loops, at least in the observed region, are connected at one foot point with IN fields we expect the chromospheric and coronal field to be more dynamic than suggested by network fields alone."
 Time series of high resohition maguctoerams in the quiet Sun can be used to investigate this further aud to revisit (og...asinvestigatedatmuchlower 2005)..," Time series of high resolution magnetograms in the quiet Sun can be used to investigate this further and to revisit \citep[e.g., as investigated at much lower spatial resolution by][]{close:etal04,close:etal05}. ."
where we have dropped the term xWyWi).,"where we have dropped the term $\propto \bnabla \Sigma_0 \times \bnabla
P_0$."
 Notice that an entropy gradient is uot required for the evolution of the perturbed potential vorticity., Notice that an entropy gradient is not required for the evolution of the perturbed potential vorticity.
"For Sy=0X,constant. equation (?7)) reduces to TVS;! ON Potential vorticity is conserved ouly iu the limit of zero stratification (πι=constant) or adiabatic perturbatious (6.9= Q).","For $S_0 = P_0 \Sigma_0^\gamma = 
constant$ , equation \ref{PVEVLIN}) ) reduces to + _0 + _0 = Potential vorticity is conserved only in the limit of zero stratification $P_0 
= constant$ ) or adiabatic perturbations $\delta S = 0$ )."
 Our goal is to understand the effects of radial stratification. but we begin by developing the linear theory of the staucard (unstratified) shearing sheet. in which the equilibrium deusity aud pressure are assumed to be spatially coustaut.," Our goal is to understand the effects of radial stratification, but we begin by developing the linear theory of the standard (unstratified) shearing sheet, in which the equilibrium density and pressure are assumed to be spatially constant."
 This will serve to establish notation and method of analysis and to highlight the changes introduced by radial stratification in the next section., This will serve to establish notation and method of analysis and to highlight the changes introduced by radial stratification in the next section.
 Our analysis follows that of Goldreich&Tremaine(1978). except for our neglect of sell-gravity., Our analysis follows that of \cite{gt78} except for our neglect of self-gravity.
 The equilibrium: consists of a uniform sheet with X=Xyconstant. P=Pyconstant. aud vy=—qt&kreg.," The equilibrium consists of a uniform sheet with $\Sigma =
\Sigma_0 = constant$, $P = P_0 = constant$, and $\bv_0 = -q\Omega x
\ey$."
" We consider nouaxisviunetrie Eulerian perturbations about this equilibrium with space-time clepencence d(/jexp(ih,jar+duy). where kG) + yQhyl(7) (ΠΕ Ayo and Ay,>0 coustant) is required to allow for a spatial Fourier decomposition of the perturbation."," We consider nonaxisymmetric Eulerian perturbations about this equilibrium with space-time dependence $\delta(t){\rm exp} (ik_x(t) x +
ik_y y)$, where k_x(t) + k_y t (with $k_{x0}$ and $k_y > 0$ constant) is required to allow for a spatial Fourier decomposition of the perturbation."
 We will refer to these perturbatious as shearing waves. or with some trepidation. but more compactly. ass/neeaces.," We will refer to these perturbations as shearing waves, or with some trepidation, but more compactly, as."
 To linear order iu the perturbation zuuplitudes. the dyvuamical equations reduce to thy—— dey deu—— Ae m0.," To linear order in the perturbation amplitudes, the dynamical equations reduce to + i k_x v_x 	+ i k_y v_y =0,"
viewed uear 6h. UT.,viewed near 6h UT.
 Badio flux rises between tle two peaks of flare emission at 0.5 and 6.2 UT., Radio flux rises between the two peaks of flare emission at 5.8h and 6.2 UT.
toward Ligh flux values at high (S-ID/(S|IT) too. because the MECS scusitivity is reduced for very soft sources by the berilliuu window. which absorbs most pliotous below ~2 keV. The 1.5-10 keV sensitivity is miaxiuun for an unabsorbed power law spectrm of ap=0.608.,"toward high flux values at high (S-H)/(S+H) too, because the MECS sensitivity is reduced for very soft sources by the berillium window, which absorbs most photons below $\sim2$ keV. The 4.5-10 keV sensitivity is maximum for an unabsorbed power law spectrum of $\alpha_E=0.6-0.8$."
 Correlations of the IIELLAS source catalog with catalogs of cosmic sources provide 26 coincidences (7 radio-loud ACN. 13 radio-quiet. ACN. 6 clusters of galaxies). sugeesting that most of the TELLAS sources ave AGN.," Correlations of the HELLAS source catalog with catalogs of cosmic sources provide 26 coincidences (7 radio-loud AGN, 13 radio-quiet AGN, 6 clusters of galaxies), suggesting that most of the HELLAS sources are AGN."
 Optical spectroscopic follow- have been performed on about 15 HELLAS error-boxes. providing 36 new identifications (Fiore et al.," Optical spectroscopic follow-ups have been performed on about 45 HELLAS error-boxes, providing 36 new identifications (Fiore et al."
 1999. La Frauca et al.," 1999, La Franca et al."
 2000 in preparation)., 2000 in preparation).
 The ACN siuuple iucludes: 1) 28 broad liue blue contin quasars ii) 2 broad liue red! continu quasars ii) LL type L.8-1.9-2 AGN, The radio-quiet AGN sample includes: i) 28 broad line blue continuum quasars ii) 2 broad line `red' continuum quasars iii) 14 type 1.8-1.9-2 AGN
"maximum specific angular momentum /7* at the inner radial boundary. and a time-averaged value of mass accretion rate onto a central object in units of Bondi accretion rate AL,/AMLp al the end of the simulation.","maximum specific angular momentum $l^{max}_a$ at the inner radial boundary, and a time-averaged value of mass accretion rate onto a central object in units of Bondi accretion rate $\MDOT_a/\MDOT_B$ at the end of the simulation."
 To keep the same resolution at s:iall τας for all cases. our runs with lower Ry. were performed with more grid points in radial direction.," To keep the same resolution at small radii for all cases, our runs with lower $R'_S$, were performed with more grid points in radial direction."
 We followed each simulation until (he quasi-stationary state was achieved. i.e.. when the time averaged mass accretion rate and torus properties settled down.," We followed each simulation until the quasi-stationary state was achieved, i.e., when the time averaged mass accretion rate and torus properties settled down."
 Figs. l.. 2..," Figs. \ref{fig:1a}, \ref{fig:1b},"
 and 3. show the mass accretion rate evolution. for Ry=107. Ry=10|l and RY=10.. respectively. for different 5.," and \ref{fig:1c} show the mass accretion rate evolution, for $R'_S=10^{-3}$ , $R'_S=10^{-4}$, and $R'_S=10^{-5}$, respectively, for different $\gamma$."
" In all cases. after an episode of the spherically svimnnmetric inflow for which Αν —1. M,/Mpg decreases and starts to fluctuates around sole (me-averaged level."," In all cases, after an episode of the spherically symmetric inflow for which $\MDOT_a/\MDOT_B$ =1, $\MDOT_a/\MDOT_B$ decreases and starts to fluctuates around some time-averaged level."
 For ο=/5/3. (he mass accretion rate evolves in a similar wav [orall Ro. ie. it stabilizes quickly and shows no strong time-variabilitv.," For $\gamma=5/3$, the mass accretion rate evolves in a similar way forall $R'_S$ , i.e., it stabilizes quickly and shows no strong time-variability."
" However. lor 5=4/3 and 1.2. M, the amplitude of the fluctuations is significant."," However, for $\gamma=4/3$ and 1.2, $\MDOT_a$ the amplitude of the fluctuations is significant."
" In particular. M, can suddenly increase bv a factor of ~5 and then also suddenly decrease."," In particular, $\MDOT_a$ can suddenly increase by a factor of $\sim 5$ and then also suddenly decrease."
 These flares are a result of mixing of high and low angular momentum matter (see next sections)., These flares are a result of mixing of high and low angular momentum matter (see next sections).
 For ο=4/3. the occasional flaresappear for all the values of RY we explored.," For $\gamma=4/3$, the occasional flaresappear for all the values of $R'_S$ we explored."
 We suppose that the flares are also typical for flows with 5=1.2 for regardless of Fy., We suppose that the flares are also typical for flows with $\gamma=1.2$ for regardless of $R'_S$.
 This strong time variabilitw for intermediate 5. is quite surprising elven we explore a relatively simple IID case.," This strong time variability for intermediate $\gamma$, is quite surprising given we explore a relatively simple HD case."
 Another surprising result is that the amplitucle ancl time seale of the variability depends on 5., Another surprising result is that the amplitude and time scale of the variability depends on $\gamma$.
 ln particular. the Mares found in models with 7=4/3 and 1.2 disappears in models with +=1.01.," In particular, the flares found in models with $\gamma$ =4/3 and 1.2 disappears in models with $\gamma=1.01$."
" Instead in these models. AZ, exhibits small aanplitude quasi-periodic oscillations (see next sections). for all 24 values."," Instead in these models, $\MDOT_a$ exhibits small amplitude quasi-periodic oscillations (see next sections), for all $R'_S$ values."
" Comparing our results for a fixed > but various 24. we find that the M, time dependence is quite insensitive to RY."," Comparing our results for a fixed $\gamma$ but various $R'_S$, we find that the $\MDOT_a$ time dependence is quite insensitive to $R'_S$."
" This parameter determines the size of the sonic surface and therefore can somewhat affect the time-averaged M, (see Table 1).", This parameter determines the size of the sonic surface and therefore can somewhat affect the time-averaged $\MDOT_a$ (see Table 1).
" To summarize our results for M,. in Fig."," To summarize our results for $\MDOT_a$, in Fig."
 4 we plot the mass accretion rate at the end of (he simulations as a function of RY for different 5., \ref{fig:6} we plot the mass accretion rate at the end of the simulations as a function of $R'_S$ for different $\gamma$.
" The final mass accretion rate is a relatively complex function of our two model parameters: for >= 5/3. M, decreases wilh increasing Ry whereas for >=1.01 it increases with increasing fy."," The final mass accretion rate is a relatively complex function of our two model parameters: for $\gamma=5/3$ , $\MDOT_a$ decreases with increasing $R'_S$ whereas for $\gamma=1.01$ it increases with increasing $R'_S$."
" For ?= 4/3. M, is not a monotonic function of RY. ie. for RG«10* it decreases with increasingRG but for Ly10 idincreases with increasing Zi."," For $\gamma=4/3$ , $\MDOT_a$ is not a monotonic function of $R'_S$, i.e., for $R'_S<10^{-3}$ it decreases with increasing$R'_S$ but for $R'_S\simless10^{-3}$ , itincreases with increasing $R'_S$ ."
" As shown by PBO03a. M, depends on the shape and size of the sonic surface."," As shown by PB03a, $\MDOT_a$ depends on the shape and size of the sonic surface."
" We interpret the complex dependence of M, on 5 and [ο as aresult ofcomplex relation between the size", We interpret the complex dependence of $\MDOT_a$ on $\gamma$ and $R'_C$ as aresult ofcomplex relation between the size
in this relation (AIeCarthyetal.2004;O'Hara2006).,"in this relation \citep{McCarthyBBPH04,OHaraMBE06}."
. The existence of voung svstems which are out of dynamical equilibrium can also broaden the observed ÀJ—T relation. but the effects may not be very pronounced (OHaraetal.2006): we are (o some extent accounting for this. because mereing halos will have greater kinetic energy per particle and (hus a higher temperature (han a relaxed halo of the same mass.," The existence of young systems which are out of dynamical equilibrium can also broaden the observed $M-T$ relation, but the effects may not be very pronounced \citep{OHaraMBE06}; we are to some extent accounting for this, because merging halos will have greater kinetic energy per particle and thus a higher temperature than a relaxed halo of the same mass."
 We have not modeled observational error. whieh will also of course add (ο any intrinsic scatter.," We have not modeled observational error, which will also of course add to any intrinsic scatter."
 Our predicted A—T relations change little il we instead use the spectroscopic temperature. as shown in Fig. reffig:sptvsm..," Our predicted $M-T$ relations change little if we instead use the spectroscopic temperature, as shown in $.$ \\ref{fig:sptvsm}."
 Also shown are ten nearby. relaxed clusters observed with XMM-Newton by Arnaudetal.(2005). and ten relaxed. clusters observed with Chandra by Vikhlinin (2006).," Also shown are ten nearby, relaxed clusters observed with XMM-Newton by \citet{ArnaudPP05} and ten relaxed clusters observed with Chandra by \citet{VikhlininKFJMMV06}."
. Both used spectroscopic temperatures. and excluded the cores (although the racial range used to determine 7 is slightly different).," Both used spectroscopic temperatures, and excluded the cores (although the radial range used to determine $T$ is slightly different)."
 These observations exhibit. considerably smaller scatter. lendiug ereclence to the idea that differing dynamical states and cooling in cores will increase the scatter in temperature at a given mass.," These observations exhibit considerably smaller scatter, lending credence to the idea that differing dynamical states and cooling in cores will increase the scatter in temperature at a given mass."
 As was the case belore. both models agree reasonably well with the observed 3—7 relation above 4 keV. but the high feedback case is a better fit in the 23 keV range.," As was the case before, both models agree reasonably well with the observed $M-T$ relation above 4 keV, but the high feedback case is a better fit in the 2–3 keV range."
 The A—T relation can be well fit by the power law where E(z)=H(z)/Hiy., The $M-T$ relation can be well fit by the power law where $E(z)=H(z)/H_0$.
 Arnaudetal.(2005) found 44=2.6940.10 and à=1.712:0.09 for their saanple. while Vikhlininetal.(2006) obtained l=2.89d0.15. a=1.58c0.11.," \citet{ArnaudPP05} found $A=2.69\pm 0.10$ and $\alpha=1.71\pm 0.09$ for their sample, while \citet{VikhlininKFJMMV06} obtained $A=2.89\pm 0.15$, $\alpha=1.58\pm 0.11$."
 We first attempted to fit this relation to the simulated halos with ordinary least-scquares regression in (he log-log plane. but (his was dominated bv the more numerous low mass halos aud undulv affected by outliers.," We first attempted to fit this relation to the simulated halos with ordinary least-squares regression in the log-log plane, but this was dominated by the more numerous low mass halos and unduly affected by outliers."
 Thus we instead adopted the following procedure: we divided the .-axis into 20 logarithmically spaced bins. calculated the median in each bin. and then found the best fit to these points.," Thus we instead adopted the following procedure: we divided the $x$ -axis into 20 logarithmically spaced bins, calculated the median in each bin, and then found the best fit to these points."
 This in effect gives more massive clusters a higher weight in the fitting: the resulting fits follow closely the median lines in the figures., This in effect gives more massive clusters a higher weight in the fitting; the resulting fits follow closely the median lines in the figures.
" As listed in Table 4.. fitting to all halos with KT,> 3keV in the zero leedback model gives A=3.12 and a=1.49x0.02."," As listed in Table \ref{tab:olsfits}, fitting to all halos with ${\rm k}T_{sp}\ge 3$ keV in the zero feedback model gives $A=3.12$ and $\alpha=1.49\pm 0.02$."
 This follows the self-similar slope of 1.5. and gives cooler clusters al a fixed mass than is observed.," This follows the self-similar slope of 1.5, and gives cooler clusters at a fixed mass than is observed."
 In (he maximum feedback case temperatures shift to higher values. and the slope becomes steeper: ο=2.56 and à=1.6240.03.," In the maximum feedback case temperatures shift to higher values, and the slope becomes steeper: $A=2.56$ and $\alpha=1.62\pm 0.03$."
 This slope agrees well with the observations. although the normalization vields slightly hotter clusters at a given mass.," This slope agrees well with the observations, although the normalization yields slightly hotter clusters at a given mass."
 The formal error we obtain lor st is small (near one percent) so in Table 4. we give the rms fractional difference of the halos from the best-fit relation: this scatter reflects well the width of the shaded regions in the figures., The formal error we obtain for $A$ is small (near one percent) so in Table \ref{tab:olsfits} we give the $rms$ fractional difference of the halos from the best-fit relation; this scatter reflects well the width of the shaded regions in the figures.
 If the amount of feedback ej varied [rom cluster to cluster. then the scatter seen would be larger.," If the amount of feedback $\epsilon_f$ varied from cluster to cluster, then the scatter seen would be larger."
of a more sophisticated mage profile model than a Gaussian nieht not have an apparent advantage for astrometric reductions.,of a more sophisticated image profile model than a Gaussian might not have an apparent advantage for astrometric reductions.
 The slight. difference in amplitude of the sub-pixel phase corrections between CTIO aud NOFS data is surprising., The slight difference in amplitude of the sub-pixel phase corrections between CTIO and NOFS data is surprising.
 It could be caused by a slightly different. observed PSF between the two sites. even for the same secing (PWIIND.," It could be caused by a slightly different, observed PSF between the two sites, even for the same seeing (FWHM)."
 Whether this is caused by differences in the instrmucut. οπήςπας or atinosphere is currently not known.," Whether this is caused by differences in the instrument, guiding or atmosphere is currently not known."
 The small systematic errors of UCAC based positions of randomly selected Uipparcos stars confinus the good. correction of UCACS3 epoch positions as function of magnitude and color. at least for the 5 to 12 magnitude range.," The small systematic errors of UCAC based positions of randomly selected Hipparcos stars confirms the good correction of UCAC3 epoch positions as function of magnitude and color, at least for the 8 to 12 magnitude range."
 With UCACS positions agrecing with IHipparcos data the magnitude equation seen in residuals with respect to Tycho-2 is an indication for such sinall systematic errors in the Tycho-2 catalog., With UCAC3 positions agreeing with Hipparcos data the magnitude equation seen in residuals with respect to Tycho-2 is an indication for such small systematic errors in the Tycho-2 catalog.
 These are likely introduced through the proper motions. thus the early epoch. ground-based data. as also indicated in the UCACS release paper.," These are likely introduced through the proper motions, thus the early epoch, ground-based data, as also indicated in the UCAC3 release paper."
 The random crrors in the observed Hipparcos position differences are even slightly sinaller than the expected. combined formal errors.," The random errors in the observed $-$ Hipparcos position differences are even slightly smaller than the expected, combined formal errors."
 For the new Tipparcos reductions the difference is oly a few percent. while for the comparison with the original Hipparcos Catalogue data the observed errors are about smaller than expected.," For the new Hipparcos reductions the difference is only a few percent, while for the comparison with the original Hipparcos Catalogue data the observed errors are about smaller than expected."
 This indicates a slightly overestimated crror iu the original Hipparcos Catalogue proper motions., This indicates a slightly overestimated error in the original Hipparcos Catalogue proper motions.
 The ‘orlnal position errors for the individual UCAC observations do mclude the formal image profile fit error. aud the conventional plate adjustineut error xopagation.," The formal position errors for the individual UCAC observations do include the formal image profile fit error, and the conventional plate adjustment error propagation."
 The weighting scheme used in this individual CCD frame least-square adjustincuts also include au estimated error contribution frou. he turbulence iu the atinosphere. scaled by the exposure time.," The weighting scheme used in this individual CCD frame least-square adjustments also include an estimated error contribution from the turbulence in the atmosphere, scaled by the exposure time."
 The mismatch between the actua PSF aud the nuage profile model can lead to an overestimation of the center position errors. particularly for stars as bright as this sample of Ilipparcos stars. which would explain the slightly sinaller than expected scatter in the Uipparcos to UCAC position differcuces.," The mismatch between the actual PSF and the image profile model can lead to an overestimation of the center position errors, particularly for stars as bright as this sample of Hipparcos stars, which would explain the slightly smaller than expected scatter in the Hipparcos to UCAC position differences."
 The exclusion. of outliers at an arbitrary Πατ of 200 mas could be another possible explanation., The exclusion of outliers at an arbitrary limit of 200 mas could be another possible explanation.
 At the faint ere of Hipparcos (lith maguitude) the Tipparcos catalog positions are of comparable precision to typical mean UCAC positions (based ou 1 images) at their about 2000 epoch., At the faint end of Hipparcos (11th magnitude) the Hipparcos catalog positions are of comparable precision to typical mean UCAC positions (based on 4 images) at their about 2000 epoch.
 The next step after UCAC. the USNO Robotic Astrometric Telescope (URAT) program (Zacharias2008) to begin iu 2010 thus will likely be capable of improving proper imotious of individual Uipparcos stars sienificautly.," The next step after UCAC, the USNO Robotic Astrometric Telescope (URAT) program \citep{urat} to begin in 2010 thus will likely be capable of improving proper motions of individual Hipparcos stars significantly."
 The eutire UCAC team is thanked for making this alksky survey a reality., The entire UCAC team is thanked for making this all-sky survey a reality.
 For more detailed information about “who is who” in the UCAC project the reader is referred to the reacline file aud UCACS release paper.," For more detailed information about “who is who"" in the UCAC project the reader is referred to the readme file and UCAC3 release paper."
 The California Institute of Technology is acknowledged for the software., The California Institute of Technology is acknowledged for the software.
 More information about this project is available at /vww.usno.navyauill/usuo/astromoetrv/.., More information about this project is available at .
X-ray spectrum. norm=EM/(Azc) . EM—fningdV. being the emission measure ol (he hot gas and d the distance to the object.,"X-ray spectrum, $norm = EM/(4 \pi d^2)$ $^{-5}$, $EM = \int n_e n_H dV$ being the emission measure of the hot gas and $d$ the distance to the object."
 In this way. the predicted X-ray spectrum and (he predicted αν in coronal lines are internally consistent.," In this way, the predicted X-ray spectrum and the predicted flux in coronal lines are internally consistent."
 The fIuxes in Table 5. correspond to what should be observed using the slits emploved for the optical observations aud may be compared directly to the reddenimg-corrected limits observed in Table 2.., The fluxes in Table \ref{table_predint} correspond to what should be observed using the slits employed for the optical observations and may be compared directly to the reddening-corrected limits observed in Table \ref{table_obsint}.
 It is immediately evident that there are severe discrepancies in the {fax limits observed and (he fluxes. predicted. bv models lor NGC! 7009. NGC 6543. andBD+30°3639.. with the observed limits lower bv [actors of 3.5. 12. and 8. respectively.," It is immediately evident that there are severe discrepancies in the flux limits observed and the fluxes predicted by models for NGC 7009, NGC 6543, and, with the observed limits lower by factors of 3.5, 12, and 8, respectively."
 For NGC 7027. there is only a discrepancy if the plasma temperature has the lowest value of the four considered in Table 4..," For NGC 7027, there is only a discrepancy if the plasma temperature has the lowest value of the four considered in Table \ref{table_7027par}."
 The predicted fluxes for eexplain immediately why (his line was never detected., The predicted fluxes for explain immediately why this line was never detected.
 The predicted intensity of this line is often fainter than bbv 2 orders of magnitude or more., The predicted intensity of this line is often fainter than by 2 orders of magnitude or more.
 As a result. it is always expected to be below our detection limits.," As a result, it is always expected to be below our detection limits."
 Comparing the observed limits from Table Ὁ with the N-rax. based. predictions [or ii Table 5.. it is abundantly clear that there is a large discrepancy.," Comparing the observed limits from Table \ref{table_obsint} with the X-ray based predictions for in Table \ref{table_predint}, it is abundantly clear that there is a large discrepancy."
 The models consistently predict stronger eenmission than is observed. by [actors of at least a few to 12.," The models consistently predict stronger emission than is observed, by factors of at least a few to 12."
 Only in NGC 7027 do the models based upon the X-ray spectrum predict emission fainter (han our limits. provided a higher plasma temperature is adopted in this object. but we recall (hat the plasma parameters are not well constrained due to the quality of its X-ray spectrum.," Only in NGC 7027 do the models based upon the X-ray spectrum predict emission fainter than our limits, provided a higher plasma temperature is adopted in this object, but we recall that the plasma parameters are not well constrained due to the quality of its X-ray spectrum."
 In the other three objects. all analyses to date have found much lower plasma temperatures. so (is mechanism may be ruled out in (hose cases.," In the other three objects, all analyses to date have found much lower plasma temperatures, so this mechanism may be ruled out in those cases."
 Considering all of the uncertainties Iron our analvsis. we suggest that the most plausible explanation for this observed discrepancy is that the gas in the hot bubble in these thiree planetary nebulae is stronglyzron-depleted.," Considering all of the uncertainties from our analysis, we suggest that the most plausible explanation for this observed discrepancy is that the gas in the hot bubble in these three planetary nebulae is strongly."
 pprovides an illustrative example., provides an illustrative example.
 Our experiments with abundance sets demonstrates that the X-ray spectrum does not meanngfullv constrain the abundances in (he N-rav-emitting, Our experiments with abundance sets demonstrates that the X-ray spectrum does not meaningfully constrain the abundances in the X-ray-emitting
Two coordinate svstenis will be used throughout this paper: Iu this section we discuss the net acceleration and total torque that arises from the cussion of radiatiou by a body of arbitrary shape.,Two coordinate systems will be used throughout this paper: In this section we discuss the net acceleration and total torque that arises from the emission of radiation by a body of arbitrary shape.
 Frou svuuuectry. spherical bodies do not exhibit torque or acceleration. and we derive here the egencral scaling of the torque and acceleration as function of the roughuess of the bocly. or its deviations from spliericity.," From symmetry, spherical bodies do not exhibit torque or acceleration, and we derive here the general scaling of the torque and acceleration as function of the roughness of the body, or its deviations from sphericity."
 We treat ouly the effect that arises from emission of radiation siuce7. have shown that there is zero secular chauge due to absorption of radiation., We treat only the effect that arises from emission of radiation since\cite{absor} have shown that there is zero secular change due to absorption of radiation.
 We asstune that the orbital time around. the sun aud its rotational period are nou conuneusurate. aud we cau therefore perform the time average by averaging over the orbit aud the spin sequentially.," We assume that the orbital time around the sun and its rotational period are non commensurate, and we can therefore perform the time average by averaging over the orbit and the spin sequentially."
 We find it more convenieut to first fix the anele of rotation of the asteroid around its axis auc average the torque applied to the asteroid οσαπιο an orbit around the sun. aud then average the torque as the asteroid revolves around its axis.," We find it more convenient to first fix the angle of rotation of the asteroid around its axis and average the torque applied to the asteroid during an orbit around the sun, and then average the torque as the asteroid revolves around its axis."
 Iu order to calculate the effect of YORP. the complete structure of the asteroid needs to be known.," In order to calculate the effect of YORP, the complete structure of the asteroid needs to be known."
 Iu this section we consider the scaling relations of YORP aud provide order of magnitude estimate of the effect., In this section we consider the scaling relations of YORP and provide order of magnitude estimate of the effect.
 The change in the spinrate. s. of a homogeneous asteroid scales according to: where ® is the solar radiation momentum flux given by Tere £.. is the solar huuimosity. © is the eccentricity. € is the speed of light. d is the orbital semi-major axis. p ds the deusity. and # is the leugth scale of the asteroid.," The change in the spinrate, $s$, of a homogeneous asteroid scales according to: where $\Phi$ is the solar radiation momentum flux given by Here $L_{\odot}$ is the solar luminosity, $e$ is the eccentricity, $c$ is the speed of light, $d$ is the orbital semi-major axis, $\rho$ is the density, and $R$ is the length scale of the asteroid."
 The eccentricity depeudeuce arises from averaging the torque over a heliocentric orbit., The eccentricity dependence arises from averaging the torque over a heliocentric orbit.
 We construct simple models to account for the asviunetryv of the asteroid., We construct simple models to account for the asymmetry of the asteroid.
 We choose à» poiuts randomly distributed on the unit sphere aud counect them to create a tesscllation of small triangles that cucloses the asteroid (based. on the Quickliull algorithm ?.. which produces 21)l| triugular facets for a giveu s).," We choose $n$ points randomly distributed on the unit sphere and connect them to create a tessellation of small triangles that encloses the asteroid (based on the Quickhull algorithm \cite{quickhull}, which produces $2n-4$ triangular facets for a given $n$ )."
 This method of construction eliminates shadowing of one facet over another., This method of construction eliminates shadowing of one facet over another.
 For this body we now calculate the radiation effects; and their scaling with ο or with the deviation of the shape of the body from a sphere.," For this body, we now calculate the radiation effects, and their scaling with $n$ or with the deviation of the shape of the body from a sphere."
 In the estimates below. we assume à»X1.," In the estimates below, we assume $n\gg 1$."
 We define the deviation of the asteroid from a sphere with the same volume by comparing the normalized difference in their surface areas., We define the deviation of the asteroid from a sphere with the same volume by comparing the normalized difference in their surface areas.
 where Lar?/3is the voliue of the asteroid iud S$ is the surface area of the asteroid., where $4\pi r^{3}/3$ is the volume of the asteroid and $S$ is the surface area of the asteroid.
 This definition of spherical deviation will be shown to be d;x»1., This definition of spherical deviation will be shown to be $d_{s}\propto n^{-1}$.
 In order to siuplifv the cerivation we will assuue that all of the facets are equilateral aud that the center of mass (CAL. hereafter) is located at the origin.," In order to simplify the derivation we will assume that all of the facets are equilateral and that the center of mass (CM, hereafter) is located at the origin."
 If we connect cach one of the facets to the CM and thereby create a tetrahedron. the relation between 0. which is the vertex angle of a tetrahedron face. aud Qz2z/n. the solid angle that the tetrahedron subteuds. can be found by making use of L'IDuiliers The area of cach facet ds the radius of the sphere that covers the asteroid.," If we connect each one of the facets to the CM and thereby create a tetrahedron, the relation between $\theta$, which is the vertex angle of a tetrahedron face, and $\Omega\approx 2\pi/n$, the solid angle that the tetrahedron subtends, can be found by making use of L'Huilier's The area of each facet where $R$ is the radius of the sphere that covers the asteroid."
 The height of the tetraledronu is: The voltune that is enclosed between the facet aud the sphere is: The total difference iun the volue between the asteroid and the sphere is roughly 25: V: The ratio between the radii is: By equating the volume of the asteroid with the volue of a sphere with radius + we have: and aud vield: Therefore. we obtain: The augle. 5. between the normal of a nou-equilateral facet and the vector joining the spheres ceuter to the facets centroid. is oforder 0. so:," The height of the tetrahedron The volume that is enclosed between the facet and the sphere is: The total difference in the volume between the asteroid and the sphere is roughly $2n\cdot V$ : The ratio between the radii is: By equating the volume of the asteroid with the volume of a sphere with radius $r$ we have: and and yield: Therefore, we obtain: The angle, $\gamma$ , between the normal of a non-equilateral facet and the vector joining the sphere's center to the facet's centroid, is oforder $\theta$ , so:"
22001: Connolly 22002).,2001; Connolly 2002).
 An additional benefit of these large databases is the presence of new and unusual objects. albeit. in. small numbers.," An additional benefit of these large databases is the presence of new and unusual objects, albeit in small numbers."
 Phe SDSS has already. ciscovered several unusual clwarl stars (Strauss 11999: Llarvis 22001) and a large proportion of the highest-redshift quasars now known (Pan 22001)., The SDSS has already discovered several unusual dwarf stars (Strauss 1999; Harris 2001) and a large proportion of the highest-redshift quasars now known (Fan 2001).
 The morphological selection involved. in. the definition of the δα targets reduces the heterogeneity of the constituent object populations compared to the SDSS but the 100k Data Release should still include a variety of interesting objects., The morphological selection involved in the definition of the 2dFGRS targets reduces the heterogeneity of the constituent object populations compared to the SDSS but the 100k Data Release should still include a variety of interesting objects.
 In this paper we focus on the population of broad line quasars and active galactic nuclei: (AGN) with recdshifts essentially bevond the limit of the numberredshift selection function for the normal galaxies that dominate the sample., In this paper we focus on the population of broad line quasars and active galactic nuclei (AGN) with redshifts essentially beyond the limit of the number–redshift selection function for the normal galaxies that dominate the sample.
 Such a census of quasars and AGN with redshifts +20.3 has a number of applications., Such a census of quasars and AGN with redshifts $z \ga 0.3$ has a number of applications.
 Firstly. many of the quasars with redshifts zzz0.5 in the 2dECGIUS will lie close in projection o lower redshift) galaxics he ‘composite’ quasar plus ealaxy imag| having been classified as non-stellar.," Firstly, many of the quasars with redshifts $z \ga 0.5$ in the 2dFGRS will lie close in projection to lower redshift galaxies – the `composite' quasar plus galaxy image having been classified as non-stellar."
". ""ποσο quasargalax,vw pairs provide the opportunity to probe the ohvsical concitions in the interstellar medium of the galaxies via. absorpticon line. studies of the background. quasars.", These quasar–galaxy pairs provide the opportunity to probe the physical conditions in the interstellar medium of the galaxies via absorption line studies of the background quasars.
 A smaller fraction of quasars. concentrated particularly at ower redshifts. z£0.5. will show evidence for the presence of host galaxies.," A smaller fraction of quasars, concentrated particularly at lower redshifts, $z \la 0.5$, will show evidence for the presence of host galaxies."
 The sub-ssumple of such objects will include quasars with some of the most. luminous hosts. known., The sub-sample of such objects will include quasars with some of the most luminous hosts known.
 Secondly. a few objects may represent new examples. of strong gravitational lensing. adding to the still very small number of multiply imaged quasars known.," Secondly, a few objects may represent new examples of strong gravitational lensing, adding to the still very small number of multiply imaged quasars known."
 Finally. as the companion 201 QSO Redshift Survey (2QZ: CCroom 22001) photometric catalogue. deliberately: includes: ον unresolved sources. the frequency ancl properties of quasars present in the 2dEC€GIU provides information relating to the completeness of the 202 as a function of image morphology and. redshift.," Finally, as the companion 2dF QSO Redshift Survey (2QZ; Croom 2001) photometric catalogue deliberately includes only unresolved sources, the frequency and properties of quasars present in the 2dFGRS provides information relating to the completeness of the 2QZ as a function of image morphology and redshift."
 Following a brief description. of the 2dEGIU data (Section 2)) the search methods emploved: are. detailed: in Section 3.., Following a brief description of the 2dFGRS data (Section \ref{section:2dfgrs}) ) the search methods employed are detailed in Section \ref{section:search}.
 Issues related to candidate selection are discussed in Section 4 and the resultant object catalogue is presented in Section 5.. together with comments on individual objects.," Issues related to candidate selection are discussed in Section \ref{section:identification} and the resultant object catalogue is presented in Section \ref{section:catalogue}, together with comments on individual objects."
 The paper concludes with a brief discussion of the potential uses of the quasar sample in Sections G and 7.., The paper concludes with a brief discussion of the potential uses of the quasar sample in Sections \ref{section:discussion} and \ref{section:conclusion}. .
 The 2dbhCRS is. formally. a spectroscopic survey of ~2.5107 apparently non-stellar objects to an isophotal magnitude limit of 6)=19.45.," The 2dFGRS is, formally, a spectroscopic survey of $\sim 2.5 \times
10^5$ apparently non-stellar objects to an isophotal magnitude limit of $\bj = 19.45$."
 Our search for quasars ancl broad. line AGN is conlined. to the first subset of 107 spectra made publicly available as the I00k Data Release (Colless 22001)., Our search for quasars and broad line AGN is confined to the first subset of $10^5$ spectra made publicly available as the 100k Data Release (Colless 2001).
 The 2dPGRS spectra are obtained. using the 2db instrument on the Anglo-Australian Telescope. (AAT) and represent ~45-minute integrations through fibres with an angular ciameter on the sky of ~2.1 aresee (see Lewisαἱ... 2002).," The 2dFGRS spectra are obtained using the 2dF instrument on the Anglo-Australian Telescope (AAT) and represent $\sim 45$ -minute integrations through fibres with an angular diameter on the sky of $\sim 2.1\,$ arcsec (see Lewis, 2002)."
 The wavelength coverageextends from approximately 3100 Πο SLOOA.., The wavelength coverageextends from approximately $3700$ to $8100$.
 Strong night-sky. emission lines. notably5571. aand6300... cllectively produce 50A--wide gaps in the spectra of all except the brightest. objects.," Strong night-sky emission lines, notably, and, effectively produce $\sim 50$ -wide gaps in the spectra of all except the brightest objects."
 Phe nominal signaltonoise (S/N) ratio for à continuum-dominatecd object at the survey limit is LO per bbin. although many of the spectra are of significantly lower quality.," The nominal signal–to–noise (S/N) ratio for a continuum-dominated object at the survey limit is $\sim 10$ per bin, although many of the spectra are of significantly lower quality."
" The majority of the targets are local galaxies. with an average redshift of £22=0.1. although there is a ~5 »er cent ""contamination by Galactic stars."," The majority of the targets are local galaxies, with an average redshift of $\langle z \rangle =
0.1$, although there is a $\sim 5$ per cent `contamination' by Galactic stars."
 Morphological classification of the photometric input catalogue was aken from the Automatic Plate Measuring (APAI) survey (Macldox 11990) of Uruitecl Winecom Schmidt Telescope (UINST) ohotographic plates., Morphological classification of the photometric input catalogue was taken from the Automatic Plate Measuring (APM) survey (Maddox 1990) of United Kingdom Schmidt Telescope (UKST) photographic plates.
 The stargalaxy separation algorithm. described in «etail by Maddox ((1990). is highly elfective when applied to isolated images. and able to select. galaxies with an elliciency the fraction of selected. objects which are galaxies) of 9T per cent.," The star–galaxy separation algorithm, described in detail by Maddox (1990), is highly effective when applied to isolated images, and able to select galaxies with an efficiency the fraction of selected objects which are galaxies) of $\sim 97$ per cent."
" Classification of ""composite"" images. where he isophotal boundaries of two or more images overlap and the APAL parameterises the resulting composite image. »esent a more dilfieult problem given the limited number of moment-based parameters provided by the APAL to describe each image."," Classification of `composite' images, where the isophotal boundaries of two or more images overlap and the APM parameterises the resulting composite image, present a more difficult problem given the limited number of moment-based parameters provided by the APM to describe each image."
 For objects with magnitudes within the range included in the 20ECGIUS essentially all close pairs of objects with angular separations < Saresee (depending on the quality of the plate material) are detected: as composite images.," For objects with magnitudes within the range included in the 2dFGRS essentially all close pairs of objects with angular separations $\la
8\,$ arcsec (depending on the quality of the plate material) are detected as composite images."
 Phe majority of such images have been eliminated from the 2dEGIU spectroscopic target catalogue through a selection according to image classification parameter ancl clirect visual inspection (Colless 22001: Section 5.4)., The majority of such images have been eliminated from the 2dFGRS spectroscopic target catalogue through a selection according to image classification parameter and direct visual inspection (Colless 2001; Section 5.4).
 However. a small fraction of star.galaxy and even star pairs are included in the 20ECGIUS target list ancl such objects are responsible for the bulk of the contamination by Calactic stars.," However, a small fraction of star–galaxy and even star–star pairs are included in the 2dFGRS target list and such objects are responsible for the bulk of the contamination by Galactic stars."
 The population of unresolved: quasars ancl ACN: is subject to exactly the same image blending that can occur with Galactic stars. but the surface density on the sky is much smaller and hence the predicted. frequency. in the 2dECGIUS survey correspondingly lower.," The population of unresolved quasars and AGN is subject to exactly the same image blending that can occur with Galactic stars, but the surface density on the sky is much smaller and hence the predicted frequency in the 2dFGRS survey correspondingly lower."
" Nonetheless. adopting a surface density of 10ος 72 for quasars with magnitudes 5,z; 19.5. simple geometric arguments preclict that the 2dFGRS I00k Data Release should contain several tens ofquasars."," Nonetheless, adopting a surface density of $10\,$ $^{-2}$ for quasars with magnitudes $b_{\rm J} \la 19.5$ , simple geometric arguments predict that the 2dFGRS 100k Data Release should contain several tens ofquasars."
diagnostics (222)...,"diagnostics \citep{burrows90,hoeflich98,gomez98}."
 As in earlier studies. we examine only a few lines and their ratios in detail.," As in earlier studies, we examine only a few lines and their ratios in detail."
 Owing to the simplicity of spectral formation in the -ray regime. exactly which lines are considered is fairly unimportant — the qualitative behaviour of any line ratio will be similar to that of one of those we show.," Owing to the simplicity of spectral formation in the $\gamma$ -ray regime, exactly which lines are considered is fairly unimportant – the qualitative behaviour of any line ratio will be similar to that of one of those we show."
" Since the strongest lines originate in the decay of ""Co. their intrinsic relative strengths are set by nuclear physics alone — thus late-time measurements of line ratios yield no information that cannot be gleaned from the energy-integrated light curve."," Since the strongest lines originate in the decay of $^{56}$ Co, their intrinsic relative strengths are set by nuclear physics alone – thus late-time measurements of line ratios yield no information that cannot be gleaned from the energy-integrated light curve."
 However. at early times. the energy sensitivity of σε. makes the intensity ratio of lines at different energies sensitive to NV. (2)..," However, at early times, the energy sensitivity of $\sigma_C$ makes the intensity ratio of lines at different energies sensitive to $N_{e}$ \citep{hoeflich98}. ."
 The left panel of Fig., The left panel of Fig.
" 5. shows the ratio #,)= //(0.847 MeV)/ (2.598 MeV).", \ref{fig:line_rat} shows the ratio $R_1 = F($ 0.847 $) / F($ 2.598 $)$.
 Both lines in this ratio are from the decay of 7 Co. We chose to consider this particular pair of lines since they are strong. unblended and well-separated in energy.," Both lines in this ratio are from the decay of $^{56}$ Co. We chose to consider this particular pair of lines since they are strong, unblended and well-separated in energy."
" Although an even more widely spaced pair of lines associated with ""Co could be achieved considering the 0.511 keV positron-annihilation line. we prefer to avoid reliance on our adopted positronium fraction."," Although an even more widely spaced pair of lines associated with $^{56}$ Co could be achieved considering the 0.511 keV positron-annihilation line, we prefer to avoid reliance on our adopted positronium fraction."
 The fy ratio in Fig., The $R_1$ ratio in Fig.
 5. has been normalised such that it tends asymptotically to 1.0 in the optically thin limit., \ref{fig:line_rat} has been normalised such that it tends asymptotically to 1.0 in the optically thin limit.
" For times up to about 70 dy. #2, discriminates between different distributions of ""Ni with zc: (see Fig. 2))."," For times up to about 70 dy, $R_1$ discriminates between different distributions of $^{56}$ Ni with $\tau_C$ (see Fig. \ref{fig:taus}) )."
 Using this ratio. spherical models in which the radioactive material is behind most opacity (Model SS) and least opacity (Model SM) can be easily discriminated from the control model.," Using this ratio, spherical models in which the radioactive material is behind most opacity (Model SS) and least opacity (Model SM) can be easily discriminated from the control model."
 Aspherical models behave similarly — in Model AO. £2) is largest if observed from the side to which the Ni blob is displaced: the effect is present but weaker in Todel AE because τε. is typically smaller.," Aspherical models behave similarly – in Model AO, $R_1$ is largest if observed from the side to which the Ni blob is displaced; the effect is present but weaker in Model AE because $\tau_C$ is typically smaller."
 As for fü. for any other pair of strong Co emission lines he flux ratio οι)F(e») will be determined by the ratio of Compton opacity at the two line energies. σε(ο)ει).," As for $R_1$, for any other pair of strong Co emission lines the flux ratio $F(\epsilon_1)/F(\epsilon_2)$ will be determined by the ratio of Compton opacity at the two line energies, $\sigma_C(\epsilon_2)/\sigma_C(\epsilon_1)$."
" We will not discuss these other potential diagnostic ratios but only comment hat the most useful ratios will always be those between lines of significantly different energy (tas in. //,).", We will not discuss these other potential diagnostic ratios but only comment that the most useful ratios will always be those between lines of significantly different energy (as in $R_1$ ).
" Thus. the ratio of the two strongest lines (£°(0.847 MeV) /2°(1.238 MeV). although easier o measure. has less diagnostic value than 7?,."," Thus, the ratio of the two strongest lines $F($ 0.847 $) / F($ 1.238 $)$ ), although easier to measure, has less diagnostic value than $R_1$."
 The second line-flux ratio shown in Fig., The second line-flux ratio shown in Fig.
" 5 is between the ""Ni 0.158 MeV line and the ""Co 0.847 MeV line. referred to as /?»."," \ref{fig:line_rat} is between the $^{56}$ Ni 0.158 MeV line and the $^{56}$ Co 0.847 MeV line, referred to as $R_{2}$."
" In he figure. it is normalised to the optically thin limit. a decreasing ‘unction of time set by the decay rates of ""Ni and """" Co. For most models. /?» gives the same relative ordering as Ry at early times."," In the figure, it is normalised to the optically thin limit, a decreasing function of time set by the decay rates of $^{56}$ Ni and $^{56}$ Co. For most models, $R_2$ gives the same relative ordering as $R_1$ at early times."
 Again. this is primarily due to the dependence of oc.," Again, this is primarily due to the energy-dependence of $\sigma_C$."
 However. the effect is weaker in A than £2) because it is opposed by an increase in the Compton continuum flux around 0.158 MeV. Disentangling these two effects would require sufficiently high quality data that the  -ray continuum level around the 0.158 MeV line can be measured.," However, the effect is weaker in $R_2$ than $R_1$ because it is opposed by an increase in the Compton continuum flux around 0.158 MeV. Disentangling these two effects would require sufficiently high quality data that the $\gamma$ -ray continuum level around the 0.158 MeV line can be measured."
" Unlike /7,. however. {ο effectively separates Model SFeR from any of the others at /.«50 dy since the 0.155 MeV line is sufficiently soft to be affected by photoabsorption."," Unlike $R_1$, however, $R_2$ effectively separates Model SFeR from any of the others at $t < 50$ dy since the $0.158$ MeV line is sufficiently soft to be affected by photoabsorption."
 Thus. in agreement with ?.. we conclude that #2 (or an equivalent ratio of one of the low-energy Ni lines to a harder Co line) provides an important diagnostic for composition.," Thus, in agreement with \cite{gomez98}, we conclude that $R_2$ (or an equivalent ratio of one of the low-energy Ni lines to a harder Co line) provides an important diagnostic for composition."
" The /?, and {ο ratios defined above provide useful diagnostics for two of the quantities that may be constrained via the -ray spectrum. the distribution of total mass and the composition of the photoabsorbing plasma."," The $R_1$ and $R_2$ ratios defined above provide useful diagnostics for two of the quantities that may be constrained via the $\gamma$ -ray spectrum, the distribution of total mass and the composition of the photoabsorbing plasma."
 Peak count rates for the relevant emission lines computed adopting the same conditions described in Section 5.1.2 are shown in Fig. 6.., Peak count rates for the relevant emission lines computed adopting the same conditions described in Section \ref{sect:lc_obs} are shown in Fig. \ref{fig:line-int}.
 For both the 0.158 and 0.847 MeV lines. the peak count rates are several thousand/MeV. The intrinsic line widths. determined by the velocity range in the ejecta. are 6 keV for the 0.158 MeV line and ~34 KeV for the 0.847 MeV line.," For both the 0.158 and 0.847 MeV lines, the peak count rates are several thousand/MeV. The intrinsic line widths, determined by the velocity range in the ejecta, are $\sim 6$ keV for the 0.158 MeV line and $\sim 34$ keV for the 0.847 MeV line."
 These widths are large enough to be moderately well-resolved by modern instruments (22)).," These widths are large enough to be moderately well-resolved by modern instruments \citealt{attie03,roques03}) )."
 For our assumed distance and integration time. one would only expect around 50 source counts in the 0.847 MeV line and fewer than 10 source counts in the 0.158 MeV line.," For our assumed distance and integration time, one would only expect around 50 source counts in the 0.847 MeV line and fewer than 10 source counts in the 0.158 MeV line."
 Owing to the combination. of fewer source photons and smaller adopted detector effective area. there are fewer counts/MeV for the 2.598 MeV line — only a handful of counts integrated over the line profile.," Owing to the combination of fewer source photons and smaller adopted detector effective area, there are fewer counts/MeV for the 2.598 MeV line – only a handful of counts integrated over the line profile."
" Thus, statistical errors caused by small number counts are very substantial in both /7?, and {ο meaning that these ratios could not contain useful information for objects at our chosen distance."," Thus, statistical errors caused by small number counts are very substantial in both $R_{1}$ and $R_{2}$, meaning that these ratios could not contain useful information for objects at our chosen distance."
 A enfold increase in source counts would be needed to make the statistical accuracy comparable to any of the differences between he models (as shown in Fig. 5»., A tenfold increase in source counts would be needed to make the statistical accuracy comparable to any of the differences between the models (as shown in Fig. \ref{fig:line_rat}) ).
 With 50 — 100 times more counts and good time coverage these ratios could contain detailed information although additional instrumental effects would further imit their value in practise., With 50 – 100 times more counts and good time coverage these ratios could contain detailed information although additional instrumental effects would further limit their value in practise.
 In light of this difficulty. in the next section we consider an alternative approach in which the spectrum is coarsely binned and yardness ratios are extracted.," In light of this difficulty, in the next section we consider an alternative approach in which the spectrum is coarsely binned and hardness ratios are extracted."
 Given the simplicity of the -ray spectrum. moderately complete information can be obtained from hardness ratios alone: this is in contrast to other wavebands where spectra are complex and not well-deseribed by photometry alone.," Given the simplicity of the $\gamma$ -ray spectrum, moderately complete information can be obtained from hardness ratios alone; this is in contrast to other wavebands where spectra are complex and not well-described by photometry alone."
 Some prospects of using relatively broad energy bands for probing the soft. continuum regions of the spectrum were discussed by ?..," Some prospects of using relatively broad energy bands for probing the soft, continuum regions of the spectrum were discussed by \cite{gomez98}."
 We extend this to consider hardness ratios involving both continuum-dominated and higher energy line-dominated bands., We extend this to consider hardness ratios involving both continuum-dominated and higher energy line-dominated bands.
" We have divided the spectrum into four energy bands. two in which the Compton continuum contributes significantly to the flux (C, and C's) and two in which strong line emission always dominates CL, and Lo."," We have divided the spectrum into four energy bands, two in which the Compton continuum contributes significantly to the flux $C_1$ and $C_2$ ) and two in which strong line emission always dominates $L_1$ and $L_2$ )."
 The bands are defined below and are indicated in Fig. 3::, The bands are defined below and are indicated in Fig. \ref{fig:spec}:
" The line ratios 77, and //» (see Section 5.2)). are affected by the change in ec: between different energies."," The line ratios $R_1$ and $R_2$ (see Section \ref{sect:line_rat}) ), are affected by the change in $\sigma_C$ between different energies."
" A similar effect can be anticipated in the hardness ratio between the two line-dominated bands CL, and L2) — such an effect is present but weak (see below).", A similar effect can be anticipated in the hardness ratio between the two line-dominated bands $L_1$ and $L_2$ ) – such an effect is present but weak (see below).
 Another effect. however. can be noted when comparing fluxes at soft energies.," Another effect, however, can be noted when comparing fluxes at soft energies."
 The strength of the Compton continuum relative to the lines increases with opacity and thus high optical depths lead to softer hardness ratios when€; or C are involved., The strength of the Compton continuum relative to the lines increases with opacity and thus high optical depths lead to softer hardness ratios when$C_1$ or $C_2$ are involved.
 This is illustrated in Fig., This is illustrated in Fig.
 7. which shows four hardness ratios, \ref{fig:hard_rat} which shows four hardness ratios
compute the spectral behaviour in the X-Ray energy range of the BeppoSAX/SWIFT satellites (i.e. 1.2 keV « E « 11 keV).,compute the spectral behaviour in the X-Ray energy range of the BeppoSAX/SWIFT satellites (i.e. $1.2$ keV $<$ E $<$ $11$ keV).
 The model predicts the source to be spectrally steady in this regime with a photon index ary=—2.68 for outbursts on timescales of days., The model predicts the source to be spectrally steady in this regime with a photon index $\alpha_{\text{xray}} = -2.68$ for outbursts on timescales of days.
" Considering shorter averaging timescales of the outburst of,, e.g. the first or last two hours, two hours around the peak in the lightcurve, the maximum derivation from a;4y (@vHE) predicted by the model is +0.05 (—0.07), which could not be measured with current experiments and thus is considered as spectrally steady in this case."," Considering shorter averaging timescales of the outburst of, e.g. the first or last two hours, two hours around the peak in the lightcurve, the maximum derivation from $\alpha_{\text{xray}}$ $(\alpha_{\text{VHE}})$ predicted by the model is $\pm 0.05$ $(-0.07)$, which could not be measured with current experiments and thus is considered as spectrally steady in this case."
 Our results clearly show that the latest observations from the VERITAS telescope for still agree with a constant (steady state) emission from a SSC model when averaged over a long observation period., Our results clearly show that the latest observations from the VERITAS telescope for still agree with a constant (steady state) emission from a SSC model when averaged over a long observation period.
 This is due to the relatively moderate variability of compared to the observation TThe variability may be well explained in the context of the self-consistent treatment of acceleration of electrons in the jet., This is due to the relatively moderate variability of compared to the observation The variability may be well explained in the context of the self-consistent treatment of acceleration of electrons in the jet.
" We are aware that an outburst of the timescale of roughly five days as measured from does not necessarily require a shock in jet model, which scales down to a few minutes depending on the SSC parameters (?),, but may also be explained as e.g. different accretion states."," We are aware that an outburst of the timescale of roughly five days as measured from does not necessarily require a shock in jet model, which scales down to a few minutes depending on the SSC parameters \citep{weidinger2010b}, but may also be explained as e.g. different accretion states."
" Nevertheless the fundamental statement remains the same: long time observation of slightly variable blazars will result in a steady state emission, while an average over a single outburst will, of course, result in a significantly different SED for the source."," Nevertheless the fundamental statement remains the same: long time observation of slightly variable blazars will result in a steady state emission, while an average over a single outburst will, of course, result in a significantly different SED for the source."
 We are not yet able to rule out different emission models or even complex geometries of the emitting region., We are not yet able to rule out different emission models or even complex geometries of the emitting region.
" But we are able to model the influence of short outbursts of a source on the SED and the lightcurves in the different energy bands The VERITAS collaboration only shows an integrated spectrum for1218+30.4,, which is due to the low flux of the source and the photon index behaviour of the combined high-states."," But we are able to model the influence of short outbursts of a source on the SED and the lightcurves in the different energy bands The VERITAS collaboration only shows an integrated spectrum for, which is due to the low flux of the source and the photon index behaviour of the combined high-states."
 This integrated spectrum does not show strong variations with regard to the known low-state observed by MAGIC., This integrated spectrum does not show strong variations with regard to the known low-state observed by MAGIC.
" Our model now predicts a clear change in the spectrum, which is indicated by the dashed line in Fig."," Our model now predicts a clear change in the spectrum, which is indicated by the dashed line in Fig."
" 1, which shows the average over one outburst with a slight, currently not detectable spectral softening in the VHE range, while the synchrotron peak in the BeppoSAX/SWIFT regime remains spectrally constant."," 1, which shows the average over one outburst with a slight, currently not detectable spectral softening in the VHE range, while the synchrotron peak in the BeppoSAX/SWIFT regime remains spectrally constant."
" This situation changes for shorter and/or stronger outbursts of an overall timescale of hours, which will result in measurable spectral evolutions in all energy regimes when considered with the presented model."," This situation changes for shorter and/or stronger outbursts of an overall timescale of hours, which will result in measurable spectral evolutions in all energy regimes when considered with the presented model."
 Furthermore the time-resolved SEDs duringa flare are comprehensible with our model., Furthermore the time-resolved SEDs during a flare are comprehensible with our model.
" Hence with better time-resolved spectra or/and better multiwavelength coverage it should be possible to prove this model, and if the model is indeed applicable it will be a good tool to investigate the whole SED during an outburst without having all energy regimes observationally"," Hence with better time-resolved spectra or/and better multiwavelength coverage it should be possible to prove this model, and if the model is indeed applicable it will be a good tool to investigate the whole SED during an outburst without having all energy regimes observationally"
moclels of the ONC. we still considered a rauge of values to test the sensitivity of these calculatious to the choice ofμη.,"models of the ONC, we still considered a range of values to test the sensitivity of these calculations to the choice of."
 The observations of circumstellar disks suggest that these disks do not typically exist far beyond ages of  6 Myr., The observations of circumstellar disks suggest that these disks do not typically exist far beyond ages of $\sim$ 6 Myr.
" For our initial models of ONC. we chose values of mye, of ο. 3. aud 6 Myr."," For our initial models of ONC, we chose values of $\tau_{disk}$ of 0, 3, and 6 Myr."
 The choice of τι=0 Myr is motivated to test the idea that stellar winds aloue cau account [or the augular momentum loss between 1 and 120 Myr., The choice of $\tau_{disk} = 0$ Myr is motivated to test the idea that stellar winds alone can account for the angular momentum loss between 1 and 120 Myr.
 Disks of 6 Myr aud 3 Myr were clioseu to explore the maximally allowed and intermediate scenarios respectively., Disks of 6 Myr and 3 Myr were chosen to explore the maximally allowed and intermediate scenarios respectively.
 The four values of used in 833.1. were also used here., The four values of used in 3.1 were also used here.
 The combination of these parameters gives us 12 different augular momentui loss rates., The combination of these parameters gives us 12 different angular momentum loss rates.
 Figure 2 shows the cumulative distribution of for the low-mass ONC stars projected forward to au age of 120 Myr for these models., Figure 2 shows the cumulative distribution of for the low-mass ONC stars projected forward to an age of 120 Myr for these models.
 Figure 2 is a erid of different models. with the rows being saturation thresholds of 1.5. 3.6. 5.1. aud 7.2 w.. aud the columns being disk lockiug lifetimes of 0. 3. aud 6 Myr.," Figure 2 is a grid of different models, with the rows being saturation thresholds of 1.8, 3.6, 5.4, and 7.2 $\omega_\odot$, and the columns being disk locking lifetimes of 0, 3, and 6 Myr."
 Table 1 shows the values returued by the Ix-8 test [or all twelve panels in Figure 2., Table 1 shows the values returned by the K-S test for all twelve panels in Figure 2.
 The distributious in the first coblumu. showing the projectious of the ONC at au age of 120 Myr witl uo disk locking. are not compatible with the observed distribution of the Pleiades at confidence.," The distributions in the first column, showing the projections of the ONC at an age of 120 Myr with no disk locking, are not compatible with the observed distribution of the Pleiades at confidence."
 Iudeed. a siguificaut fraction of the stars woule be rotating at or near breakup velocity.," Indeed, a significant fraction of the stars would be rotating at or near breakup velocity."
 The projectious shown in the panels in the right-hane columu. where all stars had disks that lasted for 6 Myr. are also excluded at the same coulideuce.," The projections shown in the panels in the right-hand column, where all stars had disks that lasted for 6 Myr, are also excluded at the same confidence."
 The results of the models lor 3 Myr disks are quite good. with Ix-5 values above the 0.1 level for al four values of«vj.," The results of the models for 3 Myr disks are quite good, with K-S values above the 0.1 level for all four values of."
.. However. to assume that all stars are locked to disks ofthe saine lifetime is too simplistie to be takeu as a realistic model. aud it does uot agree with indepeucent direct estimates of aceretion disk lifetimes.," However, to assume that all stars are locked to disks of the same lifetime is too simplistic to be taken as a realistic model, and it does not agree with independent direct estimates of accretion disk lifetimes."
 But Figure 2 clearly shows that angular momeutum loss through magnetic stellar winds alone caunot account for the observatious in young stellar clusters. even when the entire range of possible values for is explored.," But Figure 2 clearly shows that angular momentum loss through magnetic stellar winds alone cannot account for the observations in young stellar clusters, even when the entire range of possible values for is explored."
 Au adcditioual loss mechanisin is required., An additional loss mechanism is required.
Racio svnchrotron emission of high energy electrons. in the interstellar medium. (SM). indicates the presence of maenetic fields in galaxies.,Radio synchrotron emission of high energy electrons in the interstellar medium (ISM) indicates the presence of magnetic fields in galaxies.
 Rotation measures (RAL) of background. polarized sources indicate two varieties of field: a random field. which is not coherent on scales larger than the turbulence of the ISM: and a spiral ordered field which exhibits large-scale coherence (e.g. Stepanov 2008).," Rotation measures (RM) of background polarized sources indicate two varieties of field: a random field, which is not coherent on scales larger than the turbulence of the ISM; and a spiral ordered field which exhibits large-scale coherence (e.g. Stepanov 2008)."
 For a typical galaxy these fields have strengths of afew pC. In a galaxy such as 551. the coherent magnetic Ποιά is observed. to be associated with the optical spira arms (Datrikeevetal.2006).," For a typical galaxy these fields have strengths of a few $\mu$ G. In a galaxy such as 51, the coherent magnetic field is observed to be associated with the optical spiral arms \cite{PatrikeevEA06}."
. Such fields are important in star formation and the physics of cosmic ravs. and coulc also have an effect on galaxy evolution. vet. despite their importance. questions about their origin. evolution. anc structure remain largely unsolved.," Such fields are important in star formation and the physics of cosmic rays, and could also have an effect on galaxy evolution, yet, despite their importance, questions about their origin, evolution and structure remain largely unsolved."
 The Square Ixilometre Array (SILA) will help us answer questions such as these., The Square Kilometre Array (SKA) will help us answer questions such as these.
 The SKA will observe polarizec svnchrotron emission (Ciaensleretal. 2004).. and the ~ALl-Sky SKA Rotation Measure Survey” will expand RAL data sets bv five orders of magnitude. mapping the magnetic fields of galaxies in unprecedented. detail (Stepanovctal.2008:Gaensler 2006).," The SKA will observe polarized synchrotron emission \cite{GaenslerEA04}, , and the ``All-Sky SKA Rotation Measure Survey” will expand RM data sets by five orders of magnitude, mapping the magnetic fields of galaxies in unprecedented detail \cite{StepanovEA08,Gaensler06}."
. In. particular. the SILA will provide data on the evolution of galactic magnetic fields to high redshift (Gacnsler2006).. making a theoretical model [or this process very valuable.," In particular, the SKA will provide data on the evolution of galactic magnetic fields to high redshift \cite{Gaensler06}, making a theoretical model for this process very valuable."
 Here. we present such a moclel within the context of hierarchical structure formation., Here we present such a model within the context of hierarchical structure formation.
 Our model i$ based. on the observed. properties. of ealactic magnetic Lields., Our model is based on the observed properties of galactic magnetic fields.
 Observations ancl simulations suggest that the random. field is generated. by turbulence in the ESAL which is. modeled as a single-phase maenctohycrodynamic (MELD) Quid. within which magnetic ficld lines are frozen (e.g. Cho 2009).," Observations and simulations suggest that the random field is generated by turbulence in the ISM, which is modeled as a single-phase magnetohydrodynamic (MHD) fluid, within which magnetic field lines are frozen (e.g. Cho 2009)."
 Simulations have shown that in such a turbulent AHID uid. the random magnetic field energv and turbulent Iluid kinetic energy are approximately equalafter a few," Simulations have shown that in such a turbulent MHD fluid, the random magnetic field energy and turbulent fluid kinetic energy are approximately equalafter a few"
galaxies. such as the stripping of the warm gas in the outer halo o£ the inlalline ealasxies via tidal interactions with the Cluster potential.,"galaxies, such as the stripping of the warm gas in the outer halo of the infalling galaxies via tidal interactions with the cluster potential."
 Dhis process has been shown to remove the hydrogen Irom the cole disk of the galaxy. ancl thts slowly strangle its star lormation activity.," This process has been shown to remove the hydrogen from the cold disk of the galaxy, and thus slowly strangle its star formation activity."
 Possible observational canciclates Lor galaxies going through this process are red or passive spirals with negligible star lormation. found at high (Couchetal.1998). and low recshilts (vandenBergh1991). known as “anemic spirals”.," Possible observational candidates for galaxies going through this process are red or passive spirals with negligible star formation, found at high \citep{couch} and low redshifts \citep{van2}, known as ""anemic spirals""."
 ln the section we demonstrated that the SI in a ACDM previouscosmology shows variations that cannot be processaccounted for bv changes in the local density alone., In the previous section we demonstrated that the SF process in a $\Lambda$ CDM cosmology shows variations that cannot be accounted for by changes in the local density alone.
" La particular. we have shown that large-scale moculations in the fractions of rec galaxies are imnostlv clue to variations in the distributions of host) halo Inasses (as suggested by Ceccarelli,Padilla&Laimbas2008) )."," In particular, we have shown that large-scale modulations in the fractions of red galaxies are mostly due to variations in the distributions of host halo masses (as suggested by \citealt{cec}) )."
 This answers the question lor the cause of the bulk af the dependence Iound in the SI! of galaxies., This answers the question for the cause of the bulk of the dependence found in the SF of galaxies.
 However. there are still small but significant cifferenees between low local density galaxies located in cillerent large-scale environments auc their rauidomsamo miss lunetion counterparts: partictdarly lor Rai Lola). Rye(216e). Ry [CL 86e) and Bye (l.050) a PLOM Land got2.11o0) at ρε.," However, there are still small but significant differences between low local density galaxies located in different large-scale environments and their random/same mass function counterparts; particularly for $R_{H1}$ $4.51\sigma$ ), $R_{H2}$ $2.16\sigma$ ), $R_{V1}$ $1.36\sigma$ ) and $R_{V2}$ $1.05\sigma$ ) at $\rho_{LOW}$, and $R_{H2}$ $2.14\sigma$ ) at $\rho_{MID}$."
 There are different possible scenarios do explain the additional changes in galaxy colours to those coming [rom the Ins Ποσαν, There are different possible scenarios to explain the additional changes in galaxy colours to those coming from the mass function.
 In the literature proposes two one related to particular.coherent dynamics of the environments of possibilities.galaxies and the other to the age o£ the galaxy it should be noticed that it is likely that these twa quantities population:are correlated.," In particular, the literature proposes two possibilities, one related to coherent dynamics of the environments of galaxies and the other to the age of the galaxy population; it should be noticed that it is likely that these two quantities are correlated."
 Afpora/£ , $M_{\rm total}/L$ 
The N-rav spectrum of the cluster of galaxies 22199 has been studied by iiiv imstriunenuts.,The X-ray spectrum of the cluster of galaxies 2199 has been studied by many instruments.
 The detection of à soft N-rav excess in this cluster was first claumed by Bowveretal.(1998)... based oon data. zouthough they do not represent anv analvsis work.," The detection of a soft X-ray excess in this cluster was first claimed by \citet{bowyer98}, based upon data, although they do not represent any analysis work."
 Kaastractal.(1999) analyzed theDeppoSA.X data of this cluster of galaxies aud found evidence of both a soft aud a hard N-rav excess at radii larger than kkpc., \citet{kaastra99} analyzed the data of this cluster of galaxies and found evidence of both a soft and a hard X-ray excess at radii larger than kpc.
 This last analysis was based onu spatially resolved spectroscopy with data from theBeppoSAN.EUVE audROSAT wissious.," This last analysis was based on spatially resolved spectroscopy with data from the, and missions."
 However. in a recent Letter. DerghóferaudBowyer (2002)... (hereafter BB) made categorical statements that the aualvsis of I&aastra et al.," However, in a recent Letter, \citet{berghoefer02}, (hereafter BB) made categorical statements that the analysis of Kaastra et al."
 is dawed., is flawed.
 DD sav explicitly that Unfortunately. he telescope sensitivity profile used is likely o be incorrect.” that“BeppoSANXN LECS does rot detect an EUV excess when the data are analyzed correctly.” that “using a procedure etter suited to the analysis of extended sources we show that there is no excess in Abell 2199.7 and that “these findines appeared to support the (incorrect) finding of an excess in this cluster usineEUVE data”," BB say explicitly that ""Unfortunately, the telescope sensitivity profile used is likely to be incorrect,"" that LECS does not detect an EUV excess when the data are analyzed correctly,"" that ""using a procedure better suited to the analysis of extended sources we show that there is no excess in Abell 2199,"" and that ""these findings appeared to support the (incorrect) finding of an excess in this cluster using data."""
 In this Letter. we discuss the receut conclusions of BB. aud we show that BB in fact did," In this Letter, we discuss the recent conclusions of BB, and we show that BB in fact did"
and has (he fewest number of [use clusters such as [rom close pairs of stars.,and has the fewest number of false clusters such as from close pairs of stars.
 Based on artificial cluster experiments. we find that (he completeness of clusters in the nuclear region drops quickly for magnitudes Fainter (hag nn:se21.25. more (han a magnitude brighter (han for the non-nuclear region.," Based on artificial cluster experiments, we find that the completeness of clusters in the nuclear region drops quickly for magnitudes fainter than $m_V\approx 21.25$, more than a magnitude brighter than for the non-nuclear region."
 This is corroborated by the fact that the luminosity distributions for the clusters begin to flatten below these magnitudes., This is corroborated by the fact that the luminosity distributions for the clusters begin to flatten below these magnitudes.
 size measurements can also be mace using the software (Larsen 1999). as described in Section 3.1.1.," Size measurements can also be made using the software (Larsen 1999), as described in Section 3.1.1."
 We measured the size of all zz68.000 sources in our Daophot catalog usingIshape.. aad selected cluster candidates to have measurements between 0.5 ancl 10 pixels (i.e.. al least 0.5 pixels broader than the point spread function) and à S/N of at least 50.," We measured the size of all $\approx68,000$ sources in our Daophot catalog using, and selected cluster candidates to have measurements between 0.5 and 10 pixels (i.e., at least 0.5 pixels broader than the point spread function) and a S/N of at least 50."
 We then applied the same neighbor criteria as used for the C-selected cluster catalog. io eliminate remaining contaminants.," We then applied the same neighbor criteria as used for the $C$ -selected cluster catalog, to eliminate remaining contaminants."
 We find overlap between (he C ancl selected catalogs to be 280%... but that ean El to fit good clusters in more crowded regions. such as within star forming complexes along the spiral armas and in (hie nuclear region.," We find overlap between the $C$ and selected catalogs to be $\approx80$, but that can fail to fit good clusters in more crowded regions, such as within star forming complexes along the spiral arms and in the nuclear region."
 The final cataloge prepared by this approach. which we refer to as our catalog.e contains a total of 1130 cluster candidates brighter than AA223.2.," The final catalog prepared by this approach, which we refer to as our catalog, contains a total of 1130 cluster candidates brighter than $M_V\approx 23.2$."
 Approximately 180 of these candidates. mostly brighter than ma:z21.5. are in the crowded nuclear region.," Approximately 180 of these candidates, mostly brighter than $m_V\approx 21.5$, are in the crowded nuclear region."
 The SExtractor software (Bertins Arnouts 1996) is also sometimes used to select star cluster candidates., The SExtractor software (Bertins Arnouts 1996) is also sometimes used to select star cluster candidates.
 We produced an independent source list using SExtractor. and selected cluster candidates to have £A>2 pix (this is not deconvolved from the PSF. as is done with ," We produced an independent source list using SExtractor, and selected cluster candidates to have $\textit{FWHM} > 2$ pix (this is not deconvolved from the PSF, as is done with )."
Jshape)). Figure 8 shows a comparison between (the SExtractor- and cluster catalogs (and the manual catalog that will be discussed in Section 3.2.4) for a small. relatively uncrowcded region of our field.," Figure \ref{fig:man_auto} shows a comparison between the SExtractor- and Daophot-based cluster catalogs (and the manual catalog that will be discussed in Section 3.2.4) for a small, relatively uncrowded region of our field."
 We find that 36 of the 47 objects in our SExtractor catalog match objects in the Daophot catalog in this region (1.e.. v1 )).," We find that 36 of the 47 objects in our SExtractor catalog match objects in the Daophot catalog in this region (i.e., 77 )."
 SExtractor does a better job (han Daophot of ilentilving star clusters as single objects rather (han as multiple point sources. but tends (o miss some of the more dilluse objects (e.g.. two objects in the bottom left of Figure 3)). and also misses most of the very bright elusters in regions with high background. particularly in the nuclear region.," SExtractor does a better job than Daophot of identifying star clusters as single objects rather than as multiple point sources, but tends to miss some of the more diffuse objects (e.g., two objects in the bottom left of Figure \ref{fig:man_auto}) ), and also misses most of the very bright clusters in regions with high background, particularly in the nuclear region."
 We find that 275% of all 901 objects brighter than ma:=23 and outside of the nuclear region in our SExtractor catalog match those in the Daophot catalog., We find that $\approx75$ of all 901 objects brighter than $m_V=23$ and outside of the nuclear region in our SExtractor catalog match those in the Daophot catalog.
We have observed a sample of low-luminosity ETGs in the Virgo cluster using the blue peakup detector of Spitzer-IRS.,We have observed a sample of low-luminosity ETGs in the Virgo cluster using the blue peakup detector of Spitzer-IRS.
 Our aperture photometry of the spatially resolved images can be used to draw the following conclusions., Our aperture photometry of the spatially resolved images can be used to draw the following conclusions.
2010).,.
". The timescale is generally determined by F2,/c. where A, is the stellar radius because the radiative diffusion time al breakout is less than the light travel time across the star."," The timescale is generally determined by $R_*/c$, where $R_*$ is the stellar radius because the radiative diffusion time at breakout is less than the light travel time across the star."
 The largest red supergiants have aradius of ~10! em (Levesqueetal.2005) so that the longest time of the breakout event is about 1 hour.," The largest red supergiants have aradius of $\sim 10^{14}$ cm \citep{levesque05}
 so that the longest time of the breakout event is about 1 hour."
 The short timescale of the bursts makes it difficult to detect them. other than with a wide-field X-ray telescope.," The short timescale of the bursts makes it difficult to detect them, other than with a wide-field X-ray telescope."
 The breakout detection of SN 2008D with (Soderbergetal.2003) was [ortunate.," The breakout detection of SN 2008D with \citep{soderberg08}
 was fortunate."
 The situation changes if there is dense mass loss prior to the supernova (hat creates an oplically thick region., The situation changes if there is dense mass loss prior to the supernova that creates an optically thick region.
 Calculations of this case eo back to early computer simulations of supernova lieht curves., Calculations of this case go back to early computer simulations of supernova light curves.
 Model 5 of Grassbergetal.(1971). and. Model D. of include a dense circumstellar shell with radius 1015 em that determines the properties of the shock breakout., Model 5 of \cite{grassberg71} and Model B of \cite{falk77} include a dense circumstellar shell with radius $\sim10^{15}$ cm that determines the properties of the shock breakout.
 The peak Iuminositv occurs on a timescale of ~15 days in the model of Grassbergοἱal., The peak luminosity occurs on a timescale of $\sim15$ days in the model of \cite{grassberg71}.
(1971).. Grasberg&Naclvozhin(1987) considered a steady. wind of limited duration just belore the supernova., \cite{grasberg87} considered a steady wind of limited duration just before the supernova.
 Chueaiοἱal.(2004). calculated a model for SN 1994W. which allows for dense mass loss around the supernova. leading to a peak luminosity at an age of ~20 davs.," \cite{chugai04a} calculated a model for SN 1994W which allows for dense mass loss around the supernova, leading to a peak luminosity at an age of $\sim20$ days."
 A svstematie numerical study of such events has recently. been carried out by Morivaetal.(2010)., A systematic numerical study of such events has recently been carried out by \cite{moriya10}.
. Because of the high surrounding density. (he supernova can be especially bright aud this (wpe of model has been invoked for luminous supernovae.," Because of the high surrounding density, the supernova can be especially bright and this type of model has been invoked for luminous supernovae."
 In. addition to SN 1994W. (Chugaietal.2004).. application of the dense mass loss model has been made to SN 2006ev (Smith&MeCray2007).. PTF ο] 2010).. and SN 2009Kkf (Morivaetal.2010).. among others.," In addition to SN 1994W \citep{chugai04a}, application of the dense mass loss model has been made to SN 2006gy \citep{smithmccray07}, PTF 09uj \citep{ofek10}, and SN 2009kf \citep{moriya10}, among others."
 The aim here is to eive an analvlical description of the shock breakout process for the case where a radiation dominated shock propagates into the mass loss region., The aim here is to give an analytical description of the shock breakout process for the case where a radiation dominated shock propagates into the mass loss region.
 Progress on (his front has been made by Oleketal.(2010)., Progress on this front has been made by \cite{ofek10}.
. The model is presented in 2 and compared to observations in 3., The model is presented in 2 and compared to observations in 3.
" We assume that the dense mass loss can be described by a steady wind acting [ου some (nme /,,; before the supernova explosion: the extent is Ly.=Cyl. where c, is the wind velocity."," We assume that the dense mass loss can be described by a steady wind acting for some time $t_{ml}$ before the supernova explosion; the extent is $R_w=v_w t_{ml}$, where $v_w$ is the wind velocity."
 The actual case may be more complex. but is not well determined: the wind assumption has been mace in other work (Grasberg&NadvozhinAlorivaetal. 2010).," The actual case may be more complex, but is not well determined; the wind assumption has been made in other work \citep{grasberg87,ofek10,moriya10}."
". IF the mass loss is in a steady wind. the density can be specified by a density parameter. D,. scaled to a M=107.AL.vrJ| and euo10kms+ wind so that p,=5.0xLOMD,r7 in ces units."," If the mass loss is in a steady wind, the density $\rho_w=\dot M/4\pi r^2v_w\equiv Dr^{-2}$ can be specified by a density parameter, $D_*$ , scaled to a $\dot M=10^{-2}\ml$ and $v_w=10\kms$ wind so that $\rho_w=5.0\times 10^{16}D_* r^{-2}$ in cgs units."
 For (vpical supernova parameters. (he explosion drives a radiation dominated shock through the star.," For typical supernova parameters, the explosion drives a radiation dominated shock through the star."
 A radiation, A radiation
"This, however, is very likely a selection bias, because firstly, short GRBs at such high redshifts must be very luminous to be observed by GBM and secondly, short GRBs are subluminous in the optical band (?) and therefore it is difficult to obtain a redshift measurement.","This, however, is very likely a selection bias, because firstly, short GRBs at such high redshifts must be very luminous to be observed by GBM and secondly, short GRBs are subluminous in the optical band \citep{kann08} and therefore it is difficult to obtain a redshift measurement."
" In order to explain the detection rate of GRBs at high-z, ? conclude that high-z GRBs must be more common (e.g. and/or intrinsically more luminous (?) than bursts at z (butsee?).."," In order to explain the detection rate of GRBs at $z$, \citet{salva07} conclude that $z$ GRBs must be more common \citep[e.g.][]{daigne06,wang11} and/or intrinsically more luminous \citep{salva09} than bursts at $z$ \citep[but see][]{but10}."
" As already mentioned above, ? found a tight correlation between the 1-5 peak-luminosity (Lp) and in GRBs."," As already mentioned above, \citet{yonetoku04} found a tight correlation between the 1-s peak-luminosity $L_p$ ) and in GRBs."
" Assuming that the luminosity function of GRBs Epindeed evolves with redshift and that the Yonetoku relation is valid, we would also expect a positive correlation of with 7."," Assuming that the luminosity function of GRBs indeed evolves with redshift and that the Yonetoku relation is valid, we would also expect a positive correlation of with $z$."
 In Fig.10 we present vs z., In \ref{fig:epvsz} we present vs $z$.
" As was shown in Sect.3.1.1 and in Table 1,, GBM can reliably measure down to ~15 keV. The solid line indicates this redshift-corrected lower limit."," As was shown in \ref{subsubsec:instbias} and in Table \ref{tab:sim}, GBM can reliably measure down to $\sim 15$ keV. The solid line indicates this redshift-corrected lower limit."
" The Spearman’s rank correlation, using only the long GRBs, is p=0.58 with a chance probability of P=2x103."," The Spearman's rank correlation, using only the long GRBs, is $\rho=0.58$ with a chance probability of $P=2\times10^{-3}$."
" When including the short GRBs, the correlation coefficient effectively remains unchanged, whereas the chance probability increases to P=4x1073, making a correlation slightly less likely."," When including the short GRBs, the correlation coefficient effectively remains unchanged, whereas the chance probability increases to $P=4\times10^{-3}$, making a correlation slightly less likely."
" However, this correlation can be explained entirely by selection effects: GRBs do not populate the empty area in Fig."," However, this correlation can be explained entirely by selection effects: GRBs do not populate the empty area in Fig."
 10 (low and z> 1) because they simply can not be detected by GBM., \ref{fig:epvsz} (low and $z>1$ ) because they simply can not be detected by GBM.
" Even though GBM recover a low value of such GRBs, as was shown above, the detection of such events is very challenging because of the low photon fluxes of these events."," Even though GBM recover a low value of such GRBs, as was shown above, the detection of such events is very challenging because of the low photon fluxes of these events."
 As one can see in Fig.10 the lower boundary of the apparent correlation is composed of the bursts that have relatively low peak photon fluxes., As one can see in \ref{fig:epvsz} the lower boundary of the apparent correlation is composed of the bursts that have relatively low peak photon fluxes.
 This is already an indication that these events reside at the lower fluence limit for GBM to both trigger on these events., This is already an indication that these events reside at the lower fluence limit for GBM to both trigger on these events.
" Since we actually know the intrinsic parameters of the 32 GRBs, one can test up to which maximum redshift, zmax these bursts could have been detected, i.e. for which GBM would have triggered."," Since we actually know the intrinsic parameters of the 32 GRBs, one can test up to which maximum redshift, $z_{\rm{max}}$ these bursts could have been detected, i.e. for which GBM would have triggered."
 GBM has many trigger algorithms (various trigger time scales for various energy ranges)., GBM has many trigger algorithms (various trigger time scales for various energy ranges).
" For the purpose of this test, we focus on the 50 keV to 300 range which is the classical trigger energy range for a GRB and a timescale of a maximum of 4.096 s for long GRBs and 1.024 s for short GRBs."," For the purpose of this test, we focus on the $50$ keV to $300$ range which is the classical trigger energy range for a GRB and a timescale of a maximum of $4.096$ s for long GRBs and $1.024$ s for short GRBs."
" In order to shift a GRB to a higher redshift, three observables change: While it is straightforward to account for the changes of andToo, the proper treatment of the flux is more complex."," In order to shift a GRB to a higher redshift, three observables change: While it is straightforward to account for the changes of and, the proper treatment of the flux is more complex."
" RMFIT outputs the spectral parameters, including the normalization (No in ph cm? s! keV~') of the spectrum, which can be recognized as a proxy for the flux of a GRB."," RMFIT outputs the spectral parameters, including the normalization $N_0$ in ph $^{-2}$ $^{-1}$ $^{-1}$ ) of the spectrum, which can be recognized as a proxy for the flux of a GRB."
" Therefore, in order to decrease the flux when shifting the GRB to ever higher redshifts, the normalization has to be decreased accordingly."," Therefore, in order to decrease the flux when shifting the GRB to ever higher redshifts, the normalization has to be decreased accordingly."
 This was done as follows:, This was done as follows:
the de-trended light curvespaper.,the de-trended light curves.
 Plot (d) highlights a typical section of data in which the high frequency systematic is prevalent and well-removed by the method., Plot (d) highlights a typical section of data in which the high frequency systematic is prevalent and well-removed by the method.
" We compared results obtained when correcting all light curves together, with that produced by treating each separately."," We compared results obtained when correcting all light curves together, with that produced by treating each separately."
" The mod.out separated batches are most suitable since they use the same CCD and channel, which leads to more closely correlated systematics, for example, the reaction motor effects are more prevalent on some CCDs than others."," The mod.out separated batches are most suitable since they use the same CCD and channel, which leads to more closely correlated systematics, for example, the reaction motor effects are more prevalent on some CCDs than others."
" The ARC performs well in the vast majority of cases, removing systematic trends without altering intrinsic stellar variability signals."," The ARC performs well in the vast majority of cases, removing systematic trends without altering intrinsic stellar variability signals."
" One remaining effect that is not currently removed by the ARC is that of the variable guide stars, primarily the eclipsing binary."," One remaining effect that is not currently removed by the ARC is that of the variable guide stars, primarily the eclipsing binary."
 These introduce small variations in some light curves but these are not frequently occurring or similar enough to be detected by the ARC., These introduce small variations in some light curves but these are not frequently occurring or similar enough to be detected by the ARC.
" A method to detect and remove this effect is currently being devised and will be included in future work, but for this study it may be omitted without significantly reducing the quality of the results."," A method to detect and remove this effect is currently being devised and will be included in future work, but for this study it may be omitted without significantly reducing the quality of the results."
" In order to ensure that we were using the original release of data correctly, we first reproduced the calculations of B11 as exactly as possible, and compared our results to theirs."," In order to ensure that we were using the original release of data correctly, we first reproduced the calculations of B11 as exactly as possible, and compared our results to theirs."
 We found no discrepancies once the different data reduction methods had been taken into account., We found no discrepancies once the different data reduction methods had been taken into account.
 There are a variety of statistics that can be used to quantify variability., There are a variety of statistics that can be used to quantify variability.
" The C11 study use the PDC ‘dispersion’ which can be downloaded from (NStED) together with other pre-computed statistics, including the light curve median and reduced x?."," The C11 study use the PDC `dispersion' which can be downloaded from (NStED) together with other pre-computed statistics, including the light curve median and reduced $\chi^2$."
 Dispersion is defined as the 1 sigma rms scatter around the median magnitude of the light curve., Dispersion is defined as the 1 sigma rms scatter around the median magnitude of the light curve.
" Instead of dispersion, B10,11 measured the light curve 'range', which is essentially a measure of the peak-to-peak variation."," Instead of dispersion, B10,11 measured the light curve `range', which is essentially a measure of the peak-to-peak variation."
 The effect of high-frequency noise was removed either by smoothing the light curve on 10-hour timescales (B10) or by discarding the upper and lower 5 percentiles (B11)., The effect of high-frequency noise was removed either by smoothing the light curve on 10-hour timescales (B10) or by discarding the upper and lower 5 percentiles (B11).
 The choice of statistic used to study the variability is somewhat arbitrary and we confirm that this choice does not significantly alter the results., The choice of statistic used to study the variability is somewhat arbitrary and we confirm that this choice does not significantly alter the results.
" We used the empirical three-section cut of Cll to distinguish between likely dwarfs and giants based on surface gravity logg and effective temperature Τομ, rather than the simpler logg cut used by B10."," We used the empirical three-section cut of C11 to distinguish between likely dwarfs and giants based on surface gravity $\log g$ and effective temperature $T_{\rm eff}$ , rather than the simpler $\log g$ cut used by B10."
" It has since been noted that the KIC contains some misidentifications (Kochetal.2010),, but since these are not expected to be numerous enough to affect our results, and for ease of comparison to C11, we used the KIC values without modification."," It has since been noted that the KIC contains some misidentifications \citep{koc10}, but since these are not expected to be numerous enough to affect our results, and for ease of comparison to C11, we used the KIC values without modification."
" Based on the method of B11, we have chosen to use the range Ryar, between the 5th and 95th percentile, for the median normalised light curve as our variability statistic."," Based on the method of B11, we have chosen to use the range $R_{\mbox{var}}$, between the 5th and 95th percentile, for the median normalised light curve as our variability statistic."
 Dispersion and reduced y? provide an appropriate measure of variability that is believed to be primarily stochastic and Gaussian., Dispersion and reduced $\chi^2$ provide an appropriate measure of variability that is believed to be primarily stochastic and Gaussian.
 Pulsations and rotational variability do not meet these criteria and therefore a measurement based on the peak-to-peak variations in the light curve is considered more relevant., Pulsations and rotational variability do not meet these criteria and therefore a measurement based on the peak-to-peak variations in the light curve is considered more relevant.
 Selecting the 5th to 95th percentile range reduces the noise on the peak-to-peak measurement., Selecting the 5th to 95th percentile range reduces the noise on the peak-to-peak measurement.
" Rvar measurements for the Raw, PDC and ARC data are compared in Figures 2 and 3.."," $R_{\mbox{var}}$ measurements for the Raw, PDC and ARC data are compared in Figures \ref{fig:comp_var} and \ref{fig:comp_var_hist}."
" The ARC clearly removes most systematic effects, reducing the lower envelope of points to the photo noise limit, however it does not have the side effect of suppressing intermediate amplitude variability (as done by the PDC; this is apparent in the scarcity of points around Ryar10?—10* ppm in Figure 3 and the middle panel of Figure 2..)"," The ARC clearly removes most systematic effects, reducing the lower envelope of points to the photo noise limit, however it does not have the side effect of suppressing intermediate amplitude variability (as done by the PDC; this is apparent in the scarcity of points around $R_{\mbox{var}} = 10^{3} - 10^{4}$ ppm in Figure \ref{fig:comp_var_hist} and the middle panel of Figure \ref{fig:comp_var}. .)"
" Another unfortunate side effect of the PDC is the introduction of high frequency noise in some light curves, which does not occur with the ARC."," Another unfortunate side effect of the PDC is the introduction of high frequency noise in some light curves, which does not occur with the ARC."
 To compare the properties of the high and low variability stars we make a cut in Ryar based on a comparison to twice the variability level of the active Sun., To compare the properties of the high and low variability stars we make a cut in $R_{\mbox{var}}$ based on a comparison to twice the variability level of the active Sun.
" The solar Ryar value was calculated from the SOHO/VIRGO summed g-r light curves for the activeSun in the year 2000, because these provide the closest match to the bandpass"," The solar $R_{\mbox{var}}$ value was calculated from the SOHO/VIRGO summed g+r light curves for the activeSun in the year 2000, because these provide the closest match to the bandpass"
2000).,.
" is an ultrasoft X-ray transient source with characteristic temperatures of a few tens of eV. It was bright in two oobservations in 2006-2007, but was not detected in a pointed observation in 1992, implying a variation factor ROSATof 264 in the 0.2-10 keV absorbed flux."," is an ultrasoft X-ray transient source with characteristic temperatures of a few tens of eV. It was bright in two observations in 2006–2007, but was not detected in a pointed observation in 1992, implying a variation factor of $\gtrsim$ 64 in the 0.2–10 keV absorbed flux."
" It was undetected again in aSwift observation in 2011 February, implying a flux decrease by a factor of 212."," It was undetected again in a observation in 2011 February, implying a flux decrease by a factor of $\gtrsim$ 12."
 It lies toward the center of the galaxy IC 4765-f01-1504 at a redshift of 0.0353., It lies toward the center of the galaxy IC 4765-f01-1504 at a redshift of 0.0353.
" No bright optical emission lines were detected from this galaxy, making this source a good tidal disruption event candidate."," No bright optical emission lines were detected from this galaxy, making this source a good tidal disruption event candidate."
 The fits to the two sspectra using a thermal disk plus a weak hard component indicate that the accretion disk luminosity appears to follow the LcT relation and that the BH mass is around 105-106Mj., The fits to the two spectra using a thermal disk plus a weak hard component indicate that the accretion disk luminosity appears to follow the $L\propto T^4$ relation and that the BH mass is around $^5$ $^6$.
". The source showed large fast variability in both oobservations, which can be explained as due to fast variations in the mass accretion rate."," The source showed large fast variability in both observations, which can be explained as due to fast variations in the mass accretion rate."
" To further check whether this is a tidal disruption event, future long-term X-ray monitoring is necessary to see whether it follows the decay expected for a tidal disruption event."," To further check whether this is a tidal disruption event, future long-term X-ray monitoring is necessary to see whether it follows the decay expected for a tidal disruption event."
 Acknowledgments: We thank the anonymous referee for the helpful comments., Acknowledgments: We thank the anonymous referee for the helpful comments.
" We acknowledge the use of public data from theROSAT,Swift and ddata archives, and the 2XMM Serendipitous Source Catalog, constructed by the XMM-Newton Survey Science Center on behalf of ESA."," We acknowledge the use of public data from the, and data archives, and the 2XMM Serendipitous Source Catalog, constructed by the XMM-Newton Survey Science Center on behalf of ESA."
 We want to thank theSwift PI Neil Gehrels for approving our ToO request to observe the field ofJ184725., We want to thank the PI Neil Gehrels for approving our ToO request to observe the field of.
1-631724..Swift is supported at PSU by NASA contract NAS5-00136., is supported at PSU by NASA contract NAS5-00136.
" The optical spectroscopy is based on observations obtained at the Gemini Observatory which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council CONICYT (Chile), the Australian Research(Canada), Council Ministérrio da Ciénncia e Tecnologia (Brazil)(Australia), and Ministerio de Ciencia, Tecnologíaa e Innovaciónn Productiva (Argentina)."," The optical spectroscopy is based on observations obtained at the Gemini Observatory which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Ministérrio da Ciênncia e Tecnologia (Brazil) and Ministerio de Ciencia, a e Innovaciónn Productiva (Argentina)."
 The observations were carried out as part of program GS-2011A-Q-90., The observations were carried out as part of program GS-2011A-Q-90.
 SAF acknowledges funding from the Australian Research Council., SAF acknowledges funding from the Australian Research Council.
"To begin with, we calculate the thermal spectrum of a geometrically thin and optically thick accretion disk around a Kerr BH.","To begin with, we calculate the thermal spectrum of a geometrically thin and optically thick accretion disk around a Kerr BH."
" Here we have four free parameters determining the luminosity (??)): the mass of the BH M, the spin parameter a, the mass accretion rate M, and the inclination angle of the disk with respect to the distant observer i."," Here we have four free parameters determining the luminosity \ref{eq-lum}) ): the mass of the BH $M$, the spin parameter $a$, the mass accretion rate $\dot{M}$, and the inclination angle of the disk with respect to the distant observer $i$."
" However, usually M and i can be deduced from independent observations (see Section 5 for an example)."," However, usually $M$ and $i$ can be deduced from independent observations (see Section \ref{s-cf} for an example)."
" The role of the spin parameter is shown in Fig. 1,,"," The role of the spin parameter is shown in Fig. \ref{f-k-spin},"
" where we assume M=10 Mo, M=1015 g/s, and i=45°."," where we assume $M = 10$ $M_\odot$, $\dot{M} = 10^{18}$ g/s, and $i = 45^\circ$."
" In the left panel, we present the radial profile of the effective temperature and, in the right panel, the observed spectrum "," In the left panel, we present the radial profile of the effective temperature and, in the right panel, the observed spectrum $\nu L(\nu)$ ."
"For a«0, we mean that the disk is counterrotating."," For $a < 0$, we mean that the disk is counterrotating."
" Since we assume that rin=rj,,4, the spin parameter determinesvL(v). the inner radius of the disk: as a increases, ri, decreases and we find warmer matter at smaller radii."," Since we assume that $r_{\rm in} = r_{\rm _{\rm ISCO}}$, the spin parameter determines the inner radius of the disk: as $a$ increases, $r_{\rm in}$ decreases and we find warmer matter at smaller radii."
 At larger radii the effective temperature is essentially independent of the spin parameter., At larger radii the effective temperature is essentially independent of the spin parameter.
" Therefore, a higher spin parameter moves the peak of vL(v) to higher frequency and to higher values."," Therefore, a higher spin parameter moves the peak of $\nu L(\nu)$ to higher frequency and to higher values."
 Changing the BH mass while keeping M has two effects., Changing the BH mass while keeping $\dot{M}$ has two effects.
" For larger masses, the effective temperature decreases (Tος.M-W and therefore the peak of spectrum moves to lower frequency."," For larger masses, the effective temperature decreases $T \propto M^{-1/2}$ ) and therefore the peak of spectrum moves to lower frequency."
" At the same time, the size of the disk increases, increasing?) the total luminosity."," At the same time, the size of the disk increases, increasing the total luminosity."
 In the left panel of Fig., In the left panel of Fig.
" 2 we show the cases M—5, 10, 15 Mo for a—0.9."," \ref{f-k-m} we show the cases $M = 5$, 10, 15 $M_\odot$ for $a = 0.9$."
 The role of the mass accretion rate is shown in the right panel of Fig. 2.., The role of the mass accretion rate is shown in the right panel of Fig. \ref{f-k-m}.
 It is clear that a change in M only changes the effective temperature., It is clear that a change in $\dot{M}$ only changes the effective temperature.
 The viewingangle 7 determines the effective disk surface seen by the distant observer and the correction due to the Doppler boosting (for i=0° there is no Doppler, The viewingangle $i$ determines the effective disk surface seen by the distant observer and the correction due to the Doppler boosting (for $i = 0^\circ$ there is no Doppler.
boosting), In Fig.
"*.. In Fig. 3 we show the cases i=5°, 45? and 85? for à—0 and a=0.99 panel)."," \ref{f-k-i} we show the cases $i = 5^\circ$, $^\circ$ and $^\circ$ for $a = 0$ (left panel) and $a = 0.99$ (right panel)."
" Lastly, (leftwe panel)show the effect of rin (rightand rout on the shape of the spectrum."," Lastly, we show the effect of $r_{\rm in}$ and $r_{\rm out}$ on the shape of the spectrum."
 So far wehave adopted the standard assumption that the inner radius of the disk is at the ISCO and we have chosen the outer radius rout=10? M., So far wehave adopted the standard assumption that the inner radius of the disk is at the ISCO and we have chosen the outer radius $r_{\rm out} = 10^3$ $M$ .
spectrin becomes quasi-contiunuous.,spectrum becomes quasi-continuous.
 The spectrum. of. damping: rates and precession: vequencies of the free ΠΕΝ modes ip to very high order are plotted as a continuous curve in Figure 3. ogether with the spectrum of the free LFCAL aud LFGOAIR modes for comparison.," The spectrum of damping rates and precession frequencies of the free HFGM modes up to very high order are plotted as a continuous curve in Figure 3, together with the spectrum of the free LFGM and LFGMR modes for comparison."
 The dotted portion of le HEGM spectrum shown in Figure 3 shows where he radial distance between corrugations becomes smaller than ↽102r.," The dotted portion of the HFGM spectrum shown in Figure 3 shows where the radial distance between corrugations becomes smaller than $10^{-2}\,r$."
 The: correspouding: precession: requencey uz is 1075. which is 10 Iz.," The corresponding precession frequency $\omega$ is 5, which is 10 Hz."
 As Figure) 3. shows. the damping.‘ rate increases. steadily with ∙↖∙∙increasing. mode number.," As Figure 3 shows, the damping rate increases steadily with increasing mode number."
 At a~:15 (ω56.3« 107). Q has fallen to about unity (0z1.1 109).," At $m \sim
15$ $\omega \approx 6.3\ee5$ ), $Q$ has fallen to about unity $\sigma \approx -1.1\ee6$ )."
 The ae spectrum breaks at this frequency. jecause the damping of ratesthe higher-order modes ↕∐↸⊳↥⋅↸∖⋜↧↴∖↴↸∖↸∖↖↽↸∖∐⋯∪↥⋅↸∖↕⋅⋜∏≻↕≺∐↖↽↖↖↽↕↑∐≺↸∖↸⊳↥⋅↸∖⋜↧↴∖↴↕∐≼↴∙⊾⊔∩⋜↧↴∖↴↕∐∖ uode fuuctious. broadeu aud thei∙ wiuding∙∙↖ uuubers increasees.," The $\omega$ $\sigma$ spectrum breaks at this frequency, because the damping rates of the higher-order modes increase even more rapidly with decreasing $\omega$ as the mode functions broaden and their winding numbers increase."
" As the mode ummuber increases. further. aud the wey:, thatfrom‘ edabove.1range ⋅ from: the: 1downward. because2) the ∣⊀∼is‘ size⋅of~O.1 ]for disk.ape dampingamps ratern coutimucs to oerow whereas‘ the frequencyi approaches the constant critical frequency weir from above."," As the mode number increases further and the precession frequency approaches $\omega_{\rm crit}$ from above, the spectrum turns sharply downward, because the damping rate continues to grow whereas the frequency approaches the constant critical frequency $\omega_{\rm crit}$ from above."
 In the hnütwea. X the tiltfunctions of the IIFCUM modes assume the same shape JairGe) Gvluch peaks near 6=1) asstmec by the tilt1 functions-+4: of: the5 LECGM1 modes∖↴⊀↴⋅ as frou. below.," In the limit, , the tiltfunctions of the HFGM modes assume the same shape $\beta_{\rm crit}(x)$ (which peaks near $x=1$ ) assumed by the tilt functions of the LFGM modes as from below."
 Hence. Although the high-order free IIFCM modes are very strongly. damped.damped.," Hence, Although the high-order free HFGM modes are very strongly damped,."
 We have found two familics of disk warping modes in: the preseucem of ↜∙eravitoniagnetie⋅∙∙ and radiation-warping torques., We have found two families of disk warping modes in the presence of gravitomagnetic and radiation-warping torques.
 Iu ecneral. all of these modes are time-dependent.," In general, all of these modes are time-dependent."
 Iu particular. the static mode found by Bardeen Pettersou (1975) is perturbed by the ⋅− ↕⋜⊔∐⋜↧↑↕≺≻∐≓↖↖⋜∐↻⋯∶↴∙↑∪↥≺∣∏↸∖⋜⋯≺∏⋝↸∖≼⊳∪↕⊔↸∖↴∖↴↖↖↽↸∖⋜∐↘↽↕⋅↖↽↑∐⊔↸∖≓ ≺∐∖↻↸," In particular, the static mode found by Bardeen Petterson (1975) is perturbed by the radiation-warping torque and becomes weakly dent."
"∖∐≼∐∐∖∐↑∙ The low-frequency. eravitomaguctic. Hueoe (LECGMB) modes have precession. nfrequencies. the lowest .frequency allowed by ,the1 recession frequency approaches which: iu: diunensiouless units: (see ]te spectrum|⋅ turus↕↙∏ sharply parameter ‘values used ‘in the & present study. up to the critical⋅⋅ frequency: weir (80e eq. [10]."," The low-frequency gravitomagnetic ing (LFGMR) modes have precession frequencies that range from the lowest frequency allowed by the size of the disk, which in dimensionless units (see 2) is $\sim\,$ 0.1 for the parameter values used in the present study, up to the critical frequency $\omega_{\rm crit}$ (see eq. \ref{wcrit}] ]),"
 which is ~1.2&104., which is $\sim1.2\ee{4}$.
 The high-frequency eravitoniaenetic (IFCALD) moces have precession frequencies, The high-frequency gravitomagnetic (HFGM) modes have precession frequencies
 observation of NGC 6240 was performed fon 27 Alarch. 1994.," observation of NGC 6240 was performed on 27 March, 1994."
 The CIS provides an X-ray image with a circular. field of view of ~50 diameter., The GIS provides an X-ray image with a circular field of view of $\sim50'$ diameter.
 The SIS was operated. with 2-CCD mode. covering a much simaller. rectangular lield of view of ~11227.," The SIS was operated with 2-CCD mode, covering a much smaller, rectangular field of view of $\sim 11' \times 22'$."
 Iwasawa Comastri (1998) presented the result. of the SIS. data only., Iwasawa Comastri (1998) presented the result of the SIS data only.
) Here. we analyse both the SIS ancl GIS data.," Here, we analyse both the SIS and GIS data."
 In particular. the CUS data are iiseful in orcler to look for any contaniination sources in the icld around GC 6240.," In particular, the GIS data are useful in order to look for any contamination sources in the field around NGC 6240."
 We accepted the GIS and SIS data that were taken when the X-Ray Telescope. (NICE: Serlemitsos et al., We accepted the GIS and SIS data that were taken when the X-Ray Telescope (XRT: Serlemitsos et al.
 1995). axis was more than 5 above the local horizon. ancl when the geomagnetic cutoll rigidity was larger than 6 CieVfe up order to ensure a low and stable background.," 1995) axis was more than $^{\circ}$ above the local horizon, and when the geomagnetic cutoff rigidity was larger than 6 GeV/c in order to ensure a low and stable background."
 Additional screening condition that the elevation angle from the sunlit eth is ereater than 25° and 20° was applied o the SISO antl SISI data. respectively.," Additional screening condition that the elevation angle from the sunlit earth is greater than $^{\circ}$ and $^{\circ}$ was applied to the SIS0 and SIS1 data, respectively."
 As reported by Turner ct al. (, As reported by Turner et al. (
1991). the data of NGC 6240 show no significant time vartation.,"1997), the data of NGC 6240 show no significant time variation."
 The GIS image in the full energy band 0.510 keV shows a point-like source that coincides with NGC 6240 witha 1 arcmin svstematic error of the satellite attitude., The GIS image in the full energy band 0.5–10 keV shows a point-like source that coincides with NGC 6240 within $\sim1$ arcmin systematic error of the satellite attitude.
 In addition to the emission from NGC 6240. the soft. Galactic Eülfuse emission covers the entire GIS field of view.," In addition to the emission from NGC 6240, the soft Galactic diffuse emission covers the entire GIS field of view."
 The posrton of NGC 6240. (/.5)2(20.73.27.29). is in the middle of an excesseemission structure of Loop 1. which is found in theROSAT All Sky Survey (Snowcden ct al.," The position of NGC 6240, $(l,b)$ $(20.73, 27.29)$, is in the middle of an excess-emission structure of Loop I, which is found in the All Sky Survey (Snowden et al."
 1997)., 1997).
 The energy spectrum of the excess emission is very soft. (AZ~0.3 keV). and it is significant only in the soft energy band. below keV. On the other hand. at higher energies. we find an extended structure as clearly seen in the image for the 10 keV band shown in Figure 3.. in which the background consisting of non-X-rav background (NXD) ancl cosmic X-ray background (CXD) has been subtracted.," The energy spectrum of the excess emission is very soft $kT\sim0.3$ keV), and it is significant only in the soft energy band below $\sim2$ keV. On the other hand, at higher energies, we find an extended structure as clearly seen in the image for the 2.4--10 keV band shown in Figure \ref{fig:gisimage2}, in which the background consisting of non-X-ray background (NXB) and cosmic X-ray background (CXB) has been subtracted."
 “These are estimated [rom the cata obtained in the blank sky observations and data taken when the XRT is pointing at the dark (night) earth (for detail see Lkebe et al., These are estimated from the data obtained in the blank sky observations and data taken when the XRT is pointing at the dark (night) earth (for detail see Ikebe et al.
 1995: Ishisakki et al., 1995; Ishisaki et al.
 1997)., 1997).
 The background. was subtractec to derive the image., The background was subtracted to derive the image.
 In Figure 3.. notable emission is located north-east of CC 6240. and extends towards north as well as east.," In Figure \ref{fig:gisimage2}, notable emission is located north-east of NGC 6240, and extends towards north as well as east."
 Since 1e extension of the emission is much larger than that of the soft. X-ray emission. of NGC 6240. detected. with41. we consider it to be an unassociated: with GC 6240.," Since the extension of the emission is much larger than that of the soft X-ray emission of NGC 6240 detected with, we consider it to be an unassociated with NGC 6240."
 In icROSAT PSPC data. there exist severa faint SOULCOS around NCC 6240 as illustrated in Figure 3.," In the PSPC data, there exist several faint sources around NGC 6240 as illustrated in Figure 3."
 These sources. if blurred with the angular response of the GIS. seem to orm a brightness structure consistent with that observed with the GIS.," These sources, if blurred with the angular response of the GIS, seem to form a brightness structure consistent with that observed with the GIS."
 Therefore. the extended: X-ray structure is most probably due to these faint background sources.," Therefore, the extended X-ray structure is most probably due to these faint background sources."
 Since he field of view of RNTE/PCA and HIZNTIS are both ~1 degree FWHIAL. the emission from the contamination sources will contribute to theRAPE spectra significantly.," Since the field of view of /PCA and HEXTE are both $\sim$ 1 degree FWHM, the emission from the contamination sources will contribute to the spectra significantly."
 Below we examine the spectrum of NGC 6240 and the properties of these contamination sources., Below we examine the spectrum of NGC 6240 and the properties of these contamination sources.
 We construct the GIS energy spectrum of NGC 6240 from a circular region of 3 arcmin radius centered on the X-ray peak., We construct the GIS energy spectrum of NGC 6240 from a circular region of 3 arcmin radius centered on the X-ray peak.
 The spectra from the two CLS sensors. GIS-2 and CGIS-3. are summed.," The spectra from the two GIS sensors, GIS-2 and GIS-3, are summed."
 The NNB | €XD background is subtracted. as has been done for the CLS image.," The NXB + CXB background is subtracted, as has been done for the GIS image."
 The soft excess emission is not subtracted. but it is negligible above 2.0 keV. The SIS spectrum of NGC 6240 is extracted: within a circle of 3 arcmin radius from the 1-0 chip-1 and the SIS-1 chip-3.," The soft excess emission is not subtracted, but it is negligible above 2.0 keV. The SIS spectrum of NGC 6240 is extracted within a circle of 3 arcmin radius from the SIS-0 chip-1 and the SIS-1 chip-3."
 The data taken from the two chips are summed together., The data taken from the two chips are summed together.
 For the SIS background. we accumulate photons from a region where no contamination source is present on the same CCD chips with which NGC 6240 was observed.," For the SIS background, we accumulate photons from a region where no contamination source is present on the same CCD chips with which NGC 6240 was observed."
 Therefore. the soft Galactic dilfuse component is also subtracted as a part of the background.," Therefore, the soft Galactic diffuse component is also subtracted as a part of the background."
(hat. even in the very. pessimistic case where 50% of the Williamsetal.(1996). objects are mergers of (wo identical galaxies. their influence on [opo] is less than 25%.,"that, even in the very pessimistic case where $50\%$ of the \citet{W96} objects are mergers of two identical galaxies, their influence on $[\sigma_{\rm BG}^2]$ is less than $25\%$."
 For example. in a perhaps more realistic situation in which about of the objects are mergers of two galaxies with different magnitudes. the effect on [oo] would smaller than ~10% (see figure 4)).," For example, in a perhaps more realistic situation in which about of the objects are mergers of two galaxies with different magnitudes, the effect on $[\sigma_{\rm BG}^2]$ would smaller than $\sim$ (see figure \ref{mergers}) )."
 In conclusion. the effect of mergers on [oj] computed using the number counts is small and can be expected to remain below about for any reasonable scenario.," In conclusion, the effect of mergers on $[\sigma_{\rm BG}^2]$ computed using the \citet{W96} number counts is small and can be expected to remain below about for any reasonable scenario."
 We can hence proceed with our test on [oj] for the [27.8. 28.8] interval using the Williamsetal.(1996) data.," We can hence proceed with our test on $[\sigma_{\rm BG}^2]$ for the [27.8, 28.8] interval using the \citet{W96} data."
 The results for SDF-measured and »(m)-estimated [oj] are eiven in Table 9.., The results for SBF-measured and $n(m)$ -estimated $[\sigma_{\rm BG}^2]$ are given in Table \ref{t-interval}.
 The value of n(m)-estimated [o3] has been computed using and Aletealleetal.(2001) data., The value of $n(m)$ -estimated $[\sigma_{\rm BG}^2]$ has been computed using \citet{W96} and \citet{Met01} data.
 The results are given in columns 2 and 3. while column 4 lists the n(n)-ameasured [ojo] values.," The results are given in columns 2 and 3, while column 4 lists the $n(m)$ -measured $[\sigma_{\rm BG}^2]$ values."
 It can be seen that they are not compatible with the values obtained Irom the Metcalfeetal...(2001) data. as expected.," It can be seen that they are not compatible with the values obtained from the \citet{Met01} data, as expected."
 On the other hand. the n(m)-estimated [ojo] obtained from the number counts and the SDF-measured [o5] are very close.," On the other hand, the $n(m)$ -estimated $[\sigma_{\rm BG}^2]$ obtained from the \citet{W96} number counts and the SBF-measured $[\sigma_{\rm BG}^2]$ are very close."
 In order to compare them with more detail. let analyze figure 4 again.," In order to compare them with more detail, let analyze figure \ref{mergers} again."
 The n(n)-measured [oj] value for the FG606W filter and ils error interval have also been plotted in the figure (shadowed region)., The $n(m)$ -measured $[\sigma_{\rm BG}^2]$ value for the F606W filter and its error interval have also been plotted in the figure (shadowed region).
 It can be seen now that they fully coincide i£ a number of mergers about is assumed., It can be seen now that they fully coincide if a number of mergers about is assumed.
 This situation is realistic aud compatible with Metcalleetal. (2001)s claims., This situation is realistic and compatible with \citet{Met01}' 's claims.
 ο the test has been successful and shows that the SBF measurements are well-calibrated.," Summarizing, the test has been successful and shows that the SBF measurements are well-calibrated."
 In this section. SDF-measured oj; results (listed in Table 7)) and the n(Gn)-estimated Gi; values obtained from both the Williamsetal.(1996) and Metcalfeetal.(2001) data (listed in Table 3)) will be compared.," In this section, SBF-measured $\sigma_{\rm BG}^2$ results (listed in Table \ref{t-results-bg}) ) and the $n(m)$ -estimated $\sigma_{\rm BG}^2$ values obtained from both the \citet{W96} and \citet{Met01} data (listed in Table \ref{t-sigma}) ) will be compared."
 In the following. two possibilities will be considered and their consequences discussed: 1) that the Williamsetal.(1996) data represent the right differential number counts. aud ii) that the Metcalleetal.(2001) number counts are correct.," In the following, two possibilities will be considered and their consequences discussed: i) that the \citet{W96} data represent the right differential number counts, and ii) that the \citet{Met01} number counts are correct."
 The obtained SBF measurements will be used to test the validityv of these possibilities and. as a result. a final (i) will be proposed.," The obtained SBF measurements will be used to test the validity of these possibilities and, as a result, a final $n(m)$ will be proposed."
"respectively,",respectively.
 Since M21 is an Sb (vpe galaxy. the bulge contribution is important. and plavs a role in determining the rotation curve even in regions outside (he bulge.," Since M31 is an Sb type galaxy, the bulge contribution is important, and plays a role in determining the rotation curve even in regions outside the bulge."
 We use (he four-parameter dark matter halo model (de Zeeuw Plenniger 1955: Becquaert Combes 1997) with the densitv profile given by where στ+(z2/q). py is the central core density of the halo. H(q) is the core radius. p is the densitv index. ancl ¢ is the vertical-to-planar axis ratio of the halo (spherical: q = lioblate: q- < 1: prolate: g > 1).," We use the four-parameter dark matter halo model (de Zeeuw Pfenniger 1988; Becquaert Combes 1997) with the density profile given by where $ m^{2}$ $R^{2} + ({z^{2}}/{q^{2}})$, $\rho_0$ is the central core density of the halo, ${R_c}(q)$ is the core radius, $p$ is the density index, and $q$ is the vertical-to-planar axis ratio of the halo (spherical: $q$ = 1; oblate: $q$ $<$ 1; prolate: $q$ $>$ 1)."
" For a given bulge and halo densitwv profile. (he equation to be solved to obtain the vertical density distribution at any radius for any component (stars. LI 7/2) is given by eq.(4). which simplifies to: This represents three coupled. second-order. ordinary differentia] equations in ps. Pury and py, which denote the mass densities for stars. HI and II» respectively."," For a given bulge and halo density profile, the equation to be solved to obtain the vertical density distribution at any radius for any component (stars, HI $H_2$ ) is given by eq.(4), which simplifies to: This represents three coupled, second-order, ordinary differential equations in $\rho_s$, $\rho_{HI}$ and $\rho_{H_2}$ which denote the mass densities for stars, HI and $_2$ respectively."
 This problem is solved in an iterative fashion. as an mitial value problem. using fourth oxder. Iunge-Ixutta method of integration. with the following two initial conditions at the mid-plaue i.e z = 0 for each component:," This problem is solved in an iterative fashion, as an initial value problem, using fourth order, Runge-Kutta method of integration, with the following two initial conditions at the mid-plane i.e z = 0 for each component:"
"to stars, that were detected by high-precision radial velocity measurements and found to have minimum masses in the brown-dwarf domain.","to stars, that were detected by high-precision radial velocity measurements and found to have minimum masses in the brown-dwarf domain."
 Our target sample is composed of stars from the and planet surveys., Our target sample is composed of stars from the and planet surveys.
" In addition, we included a list of targets selected from the literature."," In addition, we included a list of targets selected from the literature."
" To validate our method of astrometric orbit determination, we also analysed a few comparison targets with published astrometric orbits."," To validate our method of astrometric orbit determination, we also analysed a few comparison targets with published astrometric orbits."
 This study further explores the and is in line with the foregoing works by ? and ?.., This study further explores the and is in line with the foregoing works by \cite{Halbwachs:2000rt} and \cite{Zucker:2001ve}.
 The paper is organised as follows., The paper is organised as follows.
 The brown-dwarf candidates from the and surveys are presented in Sect., The brown-dwarf candidates from the and surveys are presented in Sect.
 ?? together with 18 candidates selected from the literature., \ref{sec:cortargets} together with 18 candidates selected from the literature.
 The method of combining radial-velocity orbits and Hipparcos astrometry is described in Sect. ??.., The method of combining radial-velocity orbits and Hipparcos astrometry is described in Sect. \ref{sec:method}.
 The results are summarised in Sect., The results are summarised in Sect.
 ?? and discussed in Sect. ??.., \ref{sec:results} and discussed in Sect. \ref{sec:discussion}.
 We conclude in Sect. ??.., We conclude in Sect. \ref{sec:conclusions}.
" Our target sample consists of 33 stars, that exhibit radial velocity variations caused by a companion with minimum mass M»sini in the brown-dwarf mass range of 13—80 Mj."," Our target sample consists of 33 stars, that exhibit radial velocity variations caused by a companion with minimum mass $M_2 \sin i$ in the brown-dwarf mass range of $13 - 80\, M_J$ ."
 It contains 14 stars from the survey (?) and one star from the planet search program (?).., It contains 14 stars from the survey \citep{Udry:2000kx} and one star from the planet search program \citep{Mayor:2003cs}.
" To date, the planet survey has contributed to the discovery of more than 50 extrasolar planets (e.g. ???))."," To date, the planet survey has contributed to the discovery of more than 50 extrasolar planets (e.g. \citealt{Naef:2001nx, Mayor:2004oq, Tamuz:2008qe}) )."
" is an optical echelle spectrograph mounted on the 1.2 m Swiss Telescope located at the European Southern Observatory in La Silla, Chile."," is an optical echelle spectrograph mounted on the 1.2 m Swiss Telescope located at the European Southern Observatory in La Silla, Chile."
 A description of the instrument can be found in ? and the references therein., A description of the instrument can be found in \cite{Queloz:2000ad} and the references therein.
" With its unmatched precision, the instrument permitted the discovery of more than 70 planetary companions to date, including earth-mass planets (e.g. ????))."," With its unmatched precision, the instrument permitted the discovery of more than 70 planetary companions to date, including earth-mass planets (e.g. \citealt{Pepe:2007pi, Santos:2007hb, Mayor:2009ph, Forveille:2009jt}) )."
" The characteristics, radial velocities, and orbital solutions of the 15 and stars are presented and discussed in Sects. ??--??.."," The characteristics, radial velocities, and orbital solutions of the 15 and stars are presented and discussed in Sects. \ref{sec:stellarChar}- \ref{sec:objnotes}."
" Additionally, 18 potential brown-dwarf host stars with radial-velocity orbits were selected from the literature and are listed in Sect. ??.."," Additionally, 18 potential brown-dwarf host stars with radial-velocity orbits were selected from the literature and are listed in Sect. \ref{sec:littargets}."
" These include HD 190228, initially a planet-host star (?),, whose importance we recognised during the analysis."," These include HD 190228, initially a planet-host star \citep{Sivan:2004rm}, whose importance we recognised during the analysis."
 The identifiers and basic stellar characteristics of the stars and of HIP 103019 from the survey are listed in Table 1.., The identifiers and basic stellar characteristics of the stars and of HIP 103019 from the survey are listed in Table \ref{tab:stellarCharHIP}.
" The apparent visual magnitude V, the colourB—-V,, and the parallax « are from the new Hipparcos reduction (?), whereas the spectral type is from the original Hipparcos catalogue (?)."," The apparent visual magnitude $V$, the colour, and the parallax $\varpi$ are from the new Hipparcos reduction \citep{:2007kx}, whereas the spectral type is from the original Hipparcos catalogue \citep{Perryman:1997kx}."
 Table 2 displays the derived stellar parameters., Table \ref{tab:stellarChar} displays the derived stellar parameters.
 The uncertainty of V is below 0.002 mag., The uncertainty of $V$ is below 0.002 mag.
 The absolute magnitude My is derived from the Hipparcos magnitude and parallax., The absolute magnitude $M_V$ is derived from the Hipparcos magnitude and parallax.
" The effective temperature Teg, the surface gravity logg, and the metallicity [Fe/H] are derived from the spectroscopic analysis of high-signal-to-noise spectra with the method presented in ? and using the and lines listed in ?.."," The effective temperature $T_{\mathrm{eff}}$, the surface gravity $\log g$, and the metallicity [Fe/H] are derived from the spectroscopic analysis of high-signal-to-noise spectra with the method presented in \cite{Santos:2004zl} and using the and lines listed in \citet{Sousa:2008nx}."
 The stellar rotation parameter ysini is derived from the calibration of the or cross correlation function by ?.., The stellar rotation parameter $\nu \sin i$ is derived from the calibration of the or cross correlation function by \cite{Santos:2002ad}.
" Finally, the stellar mass of the primary M, and the age are estimated from the theoretical isochrones of ? and a Bayesian estimation method described in?cgi-bin/param.."," Finally, the stellar mass of the primary $M_1$ and the age are estimated from the theoretical isochrones of \cite{Girardi:2000bf} and a Bayesian estimation method described in \cite{da-Silva:2006eu}."
 Stellar ages of main-sequence dwarfs are usually not well constrained and we quote the 1-o confidence interval., Stellar ages of main-sequence dwarfs are usually not well constrained and we quote the $\sigma$ confidence interval.
" The spectroscopic analysis was problematic for HIP 103019, which maybe due to its faintness and late spectral type."," The spectroscopic analysis was problematic for HIP 103019, which maybe due to its faintness and late spectral type."
 Errors on the parameters of this star are possibly underestimated., Errors on the parameters of this star are possibly underestimated.
 The stellar characteristics of the targets from the literature can be found in the respective references given in Sect. ??.., The stellar characteristics of the targets from the literature can be found in the respective references given in Sect. \ref{sec:littargets}.
 Optical high-resolution spectra of the stars in Table 2 were collected with the and spectrographs over time spans extending to 11 years., Optical high-resolution spectra of the stars in Table \ref{tab:stellarChar} were collected with the and spectrographs over time spans extending to 11 years.
" Radial velocities are estimated from the cross-correlations of the extracted stellar spectra with numerical templates, which depend on the target's spectral type."," Radial velocities are estimated from the cross-correlations of the extracted stellar spectra with numerical templates, which depend on the target's spectral type."
" DuringCORALIE observations, a reference Thorium-Argon spectrum is recorded simultaneously with the stellar spectrum and is used to measure and correct for residual zero-point drifts (?).."," During observations, a reference Thorium-Argon spectrum is recorded simultaneously with the stellar spectrum and is used to measure and correct for residual zero-point drifts \citep{Baranne:1996rc}."
 Photon noise limits the obtained precision to typically 2-4 ms! per epoch., Photon noise limits the obtained precision to typically 2-4 $^{-1}$ per epoch.
" To account for systematic drifts, an externalerror of 5 ms~! is quadratically added to radial-velocity uncertainties before performing the period search and the model adjustment."," To account for systematic drifts, an externalerror of 5 $^{-1}$ is quadratically added to radial-velocity uncertainties before performing the period search and the model adjustment."
" The abbreviations C98 andC07 refer to the instrument before and after its upgrade in 2007, respectively (?).."," The abbreviations C98 andC07 refer to the instrument before and after its upgrade in 2007, respectively \citep{Segransan:2010xr}. ."
 The two parameters ycogand yco; are introduced to account for the possibly differing velocity offsets of these instrument configurations., The two parameters $\gamma_{\mathrm{C98}}$and $\gamma_{\mathrm{C07}}$ are introduced to account for the possibly differing velocity offsets of these instrument configurations.
A typical result. of numerical integration of the time-dependent equations is shown in Fig.,A typical result of numerical integration of the time-dependent equations is shown in Fig.
 3., 3.
" The integration is started. close to the trivial solution 5;=Mi;0 by introducing small positive values of 2,, and AL. only.", The integration is started close to the trivial solution $R_{ij}=M_{ij}=0$ by introducing small positive values of $R_{xx}$ and $M_{xx}$ only.
 The system approaches the stable fixed. point. representing maenetized turbulence., The system approaches the stable fixed point representing magnetized turbulence.
 If one seeks a non-trivial fixed point representing a state of steady. hwerocwvnamic turbulence in rotating plane Couette flow. one obtains the solution which is the generalization of equation (34)) for the non-rotating case.," If one seeks a non-trivial fixed point representing a state of steady hydrodynamic turbulence in rotating plane Couette flow, one obtains the solution with which is the generalization of equation \ref{pcf}) ) for the non-rotating case."
" This solution is only meaningful. however. if Ry;"" is positive semi-definite. and this requires as illustrated in Fig."," This solution is only meaningful, however, if $R_{ij}$ is positive semi-definite, and this requires as illustrated in Fig."
 4. (, 4. (
In the case of equality. the solution is trivial.),"In the case of equality, the solution is trivial.)"
 Where it exists. this solution appears to be stable.," Where it exists, this solution appears to be stable."
 For the Ixeplerian case Ro=3/4 relevant to a standard aeeretion disc. steady hyvdrodynamic turbulence is possible only when C/CS/3.," For the Keplerian case ${\rm Ro}=3/4$ relevant to a standard accretion disc, steady hydrodynamic turbulence is possible only when $C_2/C_1>8/3$."
 Although this is possible in principle. such a value appears improbable a priori.," Although this is possible in principle, such a value appears improbable a priori."
" When CofC€,=1. for example. steacy turbulence is possible [or 0.145«RotLS."," When $C_2/C_1=1$, for example, steady turbulence is possible for $-0.145<{\rm Ro}^{-1}<1.145$."
" For rotation laws of the [orm Oxr"". this requires q2 1.746."," For rotation laws of the form $\Omega\propto r^{-q}$, this requires $q>1.746$ ."
 Numerical integration of the time-dependent: svstenm confirms that. when the system is perturbed. from. £3; 0. it tends towards the stable solution where this exists.," Numerical integration of the time-dependent system confirms that, when the system is perturbed from $R_{ij}=0$ , it tends towards the stable solution where this exists."
 Otherwise it exhibits decaving cpievclic oscillations about Rij=., Otherwise it exhibits decaying epicyclic oscillations about $R_{ij}=0$.
 The introduction of a magnetic perturbation allows steaclvy ALD. turbulence to developin a rotating shear How. even when Ravleigh’s stability criterion is amply satisfied. as in the Keplerian case Ro=3/4 (Fig.," The introduction of a magnetic perturbation allows steady MHD turbulence to developin a rotating shear flow, even when Rayleigh's stability criterion is amply satisfied, as in the Keplerian case ${\rm
 Ro}=3/4$ (Fig."
 5)., 5).
 Phis occurs through a non-linear magnetorotational instability (non-linear because there is no imposed magnetic lux) and the magnetic field is sustained through a non-linear dvnamo process., This occurs through a non-linear magnetorotational instability (non-linear because there is no imposed magnetic flux) and the magnetic field is sustained through a non-linear dynamo process.
 The mocel therefore reproduces in broad terms the finding of numerical studies of the non-linear evolution of the magnetorotational instability (ος. Lawley ct al., The model therefore reproduces in broad terms the finding of numerical studies of the non-linear evolution of the magnetorotational instability (e.g. Hawley et al.
 1995)., 1995).
 The assumption of an incompressible fluid. is useful in reducing the problem to a minimal. although still formidable. complexity.," The assumption of an incompressible fluid is useful in reducing the problem to a minimal, although still formidable, complexity."
 “Phere are. two reasons why the model may not. be suitable for immediate application to a compressible Iluid., There are two reasons why the model may not be suitable for immediate application to a compressible fluid.
 First. the mean density may be uniform. especially if the mean velocity. field has a non-zero divergence.," First, the mean density may be non-uniform, especially if the mean velocity field has a non-zero divergence."
 Second. the turbulence may. be essentially transonic. changing the character of the motions and Leading to greatly. enhanced. dissipation through shock formation.," Second, the turbulence may be essentially transonic, changing the character of the motions and leading to greatly enhanced dissipation through shock formation."
 The second case is bevond the scope of this paper and is unlikely to be important for the magnetorotational instability. at least in the part of an accretion dise where most of the mass resides.," The second case is beyond the scope of this paper and is unlikely to be important for the magnetorotational instability, at least in the part of an accretion disc where most of the mass resides."
 The first case. however. is important in a number of applications and can be treated in a simple wav.," The first case, however, is important in a number of applications and can be treated in a simple way."
 1n a compressible [uid it is more appropriate to celine the mean Ievnolds and Maxwell tensors as (DL which have the dimensions of stress and differ from the earlier definitions by a factor of the density., In a compressible fluid it is more appropriate to define the mean Reynolds and Maxwell tensors as and which have the dimensions of stress and differ from the earlier definitions by a factor of the density.
 Lt is convenient not to include the magnetic pressure perturbation in. AJ;;. but to regard the mean magnetic stress as dM;iM.," It is convenient not to include the magnetic pressure perturbation in $M_{ij}$ , but to regard the mean magnetic stress as $M_{ij}-{\textstyle{{1}\over{2}}}M\delta_{ij}$."
 The first issue to consider is whether a divergence of the mean velocity Geld. Όρη. should. alleet the evolution of A; or M;j.," The first issue to consider is whether a divergence of the mean velocity field, $\partial_iu_i$, should affect the evolution of $R_{ij}$ or $M_{ij}$ ."
 Εις velocity. divergence does not appear in the equation of motion. but does appear in the equation of nis conservation and the induction equation in the forms Opomep;u; and OD;=D;;u; respectively.," This velocity divergence does not appear in the equation of motion, but does appear in the equation of mass conservation and the induction equation in the forms $\partial_t\rho=\cdots-\rho\partial_iu_i$ and $\partial_tB_i=\cdots-B_i\partial_ju_j$ respectively."
" This motivates the addition of terms in the forms 0,0;[jus and OAL;2o2M;Oni, respectively.", This motivates the addition of terms in the forms $\partial_tR_{ij}=\cdots-R_{ij}\partial_ku_k$ and $\partial_tM_{ij}=\cdots-2M_{ij}\partial_ku_k$ respectively.
 The second issue concerns the definition of the vertical length-scale L in a stratified. disc. and the effect of. the stratification on the vertical profile of the stress.," The second issue concerns the definition of the vertical length-scale $L$ in a stratified disc, and the effect of the stratification on the vertical profile of the stress."
 In. the incompressible system. the turbulence is homogeneous. ancl the stress independent of 2., In the incompressible system the turbulence is homogeneous and the stress independent of $z$.
 Numerical simulations of stratified. isothermal aceretion disces (CMiller Stone 20 suggest that the stress is also stratified. being roughly. proportional to the density. (or pressure). ancl in many applications it is analytically convenient to assume that the stress scales with the pressure.," Numerical simulations of stratified, isothermal accretion discs (Miller Stone 2000) suggest that the stress is also stratified, being roughly proportional to the density (or pressure), and in many applications it is analytically convenient to assume that the stress scales with the pressure."
 This can be achieved in a natural wav by identifving the vertical length-scale as L= c.f. where e=(pp 7ods the isothermal sound. speed l rper 1ο vertical oscillation frequency. which measures the curvature of the gravitational potential and is equal to Qin a Ixeplerian disc.," This can be achieved in a natural way by identifying the vertical length-scale as $L=c_{\rm s}/\Omega_z$ , where $c_{\rm s}=(p/\rho)^{1/2}$ is the isothermal sound speed and $\Omega_z$ the vertical oscillation frequency, which measures the curvature of the gravitational potential and is equal to $\Omega$ in a Keplerian disc."
 In an isothermal disc £L is then equal to the Gaussian scale-height. 441., In an isothermal disc $L$ is then equal to the Gaussian scale-height $H$.
 Accordinglv. the system.of equations. recommencled[or a compressible accretion How. is as follows. written in an inertial frame of reference.," Accordingly, the systemof equations recommendedfor a compressible accretion flow is as follows, written in an inertial frame of reference."
 “Phe equation. of massconservation. The equation of motion.," The equation of massconservation, The equation of motion,"
turnoll and the surface rotation rate.,turnoff and the surface rotation rate.
 An extrapolation of the angular momentum loss model for Pop I stars (Sills.Pinsonneault&Terndrup1999]. yields ρε=1 Kuss ! at the turnoll aud a total angular momentum ofm 5 x10!4i ο en?L7 1⋅∢∙∩⋯↥↽≻⋜⋃⋅↩≺⇂∖∖↽∐∐−≻≍∐≻↓∖∑≟∢∙⋯−⊳∖↓↥∩↕⋅⋜↕⊳∖↕⋜∐⋅↓⋅∩⋜≹⋃∐∑≟ ⋅ i .κ > . ⋅ at the current observational limit of £ ku ↓ Guid is comparable to the angular momeutum of a solar model which rotates as a solid body).," An extrapolation of the angular momentum loss model for Pop I stars \citep{SPT99} yields $v_{surf} = 1$ km $^{-1}$ at the turnoff and a total angular momentum of 5 $\times 10^{47}$ g $^2$ $^{-1}$, compared with 2 $\times 10^{48}$ g $^2$ $^{-1}$ for a star rotating at the current observational limit of 4 km $^{-1}$ (and is comparable to the angular momentum of a solar model which rotates as a solid body)."
 We will therefore consider au initial angular momentuui of D-5x.LO ο em?2 | as our base case and i2x4107κ @ em?2 1 as our limiting case. with augular iuomentunm distributed as a function of mass the same as in a solid body rotator.," We will therefore consider an initial angular momentum of $5
\times 10^{47}$ g $^2$ $^{-1}$ as our base case and $2 \times
10^{48}$ g $^2$ $^{-1}$ as our limiting case, with angular momentum distributed as a function of mass the same as in a solid body rotator."
 Iu main sequence stars there is efficient angular momentum loss [rom a magnetic wind., In main sequence stars there is efficient angular momentum loss from a magnetic wind.
 Because of the low predicted surface rotation rates. the amount of angular momentum loss [rom a magnetic wind is expected to be sinall ou the elaut branch.," Because of the low predicted surface rotation rates, the amount of angular momentum loss from a magnetic wind is expected to be small on the giant branch."
 However. large amounts of mass loss ou tlie first ascent ejant brauch are required to explain the distribution of horizoutal branch star masses Zit 1990).. aud the matter at the surface will carry away at least its own local angular© momentum.," However, large amounts of mass loss on the first ascent giant branch are required to explain the distribution of horizontal branch star masses \citep{R73,LDZ90}, and the matter at the surface will carry away at least its own local angular momentum."
 Mass and angularo momentum loss was therefore calculated as follows., Mass and angular momentum loss was therefore calculated as follows.
 Nass loss on the giant branch was calculated using Reimers” formulation (Reimers1975):: where L is the total luminosity. g is the surface gravity. aud 2 is the radius of the star.," Mass loss on the giant branch was calculated using Reimers' formulation \citep{R75}: : where $L$ is the total luminosity, $g$ is the surface gravity, and $R$ is the radius of the star."
 The constant à can be varied to allow for cifferent amounts of mass to be lost over the entire giant brauch evolution., The constant $\alpha$ can be varied to allow for different amounts of mass to be lost over the entire giant branch evolution.
 The mass loss was assumed to begin at the point of inaxiuiunm convection zone depth in mass (see section 2.1)., The mass loss was assumed to begin at the point of maximum convection zone depth in mass (see section 2.4).
 The loss of mass was taken iuto account in the evolution of the star., The loss of mass was taken into account in the evolution of the star.
 We calculated the angular morientuin lost [rom the star as it loses mass. following the method used in Pinsouneault.Delivaunis&Demarque(1991).," We calculated the angular momentum lost from the star as it loses mass, following the method used in \cite{PDD}."
. Rotation is not included iu the structural evolution of the star. and rotational mixing is not considered in these models.," Rotation is not included in the structural evolution of the star, and rotational mixing is not considered in these models."
 As mass is lost from. the star. the fractional angular momentum loss in a giveu timestep is where AAS is the total mass lost duriug the time step. /? is the stellar radius. aud v is the rotation rate at the surface of the star.," As mass is lost from the star, the fractional angular momentum loss in a given timestep is where $\Delta M$ is the total mass lost during the time step, $R$ is the stellar radius, and $\omega$ is the rotation rate at the surface of the star."
 Models with different mass loss rates will therefore experieuce dillereut amounts of angular momentunm loss., Models with different mass loss rates will therefore experience different amounts of angular momentum loss.
 The Reimers” mass loss coustant a was varied so that different amounts of mass. up to almost all of the inass outside the heliuu buruiug core. were lost [rom the initially 0.8 ML. inodel.," The Reimers' mass loss constant $\alpha$ was varied so that different amounts of mass, up to almost all of the mass outside the helium burning core, were lost from the initially 0.8 $_{\odot}$ model."
 For each oL the resulting combinations of core mass (takenfrom the model at the tip of the giant. branch), For each of the resulting combinations of core mass (takenfrom the model at the tip of the giant branch)
"Iu the secoud subcase. recall that IX—54ν55/60(04.52)) byProposition 5.. where 54 is the covering involutiou of the double cover S$,>Ti of an Euriques surface aud 5» Is the covering involution of some other (possibly ramified)double cover So> T».","In the second subcase, recall that $X=S_1\times
S_2/\langle(\gamma_1,\gamma_2)\rangle$ byProposition \ref{Enriques}, , where $\gamma_1$ is the covering involution of the double cover $S_1\rightarrow T_1$ of an Enriques surface and $\gamma_2$ is the covering involution of some other (possibly ramified)double cover $S_2\rightarrow T_2$ ."
error circle marked.,error circle marked.
 Anemission-colour diagram similar to ligure 2 was produced. for this field., An diagram similar to Figure 2 was produced for this field.
 Onlv one Be star is found within the X-ray error circle. lying only 3 arcseconds from the X-ray position (marked as object 1 in Fig 0).," Only one Be star is found within the X-ray error circle, lying only 3 arcseconds from the X-ray position (marked as object 1 in Fig 1(b)."
 A second. Be star (Object 2 in the figure) lies 25 areseconds from the X-ray. position., A second Be star (Object 2 in the figure) lies 25 arcseconds from the X-ray position.
 Based on the positional coincidence. Object 1 is most probably the counterpart to the N-rav pulsar. although Object 2 cannot be wholly dismissed at this stage.," Based on the positional coincidence, Object 1 is most probably the counterpart to the X-ray pulsar, although Object 2 cannot be wholly dismissed at this stage."
 This source was first detected: as a 91.12. second. pulsar inANTE observations (Corbet ct al., This source was first detected as a 91.12 second pulsar in observations (Corbet et al.
 1998a) although was initially confused with the nearby 46 second. pulsar INCL J0053.8-7226 (Buckles et al., 1998a) although was initially confused with the nearby 46 second pulsar 1WGA J0053.8-7226 (Buckley et al.
 1998)., 1998).
" Further observations with revealed two pulsars in the field with an approxiniate 2 to 1 ralio in periocls. the 91 second period belonging lo the new source. ΑΝ. J0051-722. whilst, observations with ROSAT reduced. the positional uncertainty to 10 areseconds."," Further observations with revealed two pulsars in the field with an approximate 2 to 1 ratio in periods, the 91 second period belonging to the new source, AX J0051-722, whilst observations with ROSAT reduced the positional uncertainty to 10 arcseconds."
 We performed. spectroscopic observations of the brightest object in this crror cirele on 1998 February 3 (sce finder chart in Figure. 1(c))., We performed spectroscopic observations of the brightest object in this error circle on 1998 February 3 (see finder chart in Figure 1(c)).
 The spectrum (Figure 3(b)) shows the La line strongly in emission. with an equivalent width of -22A.," The spectrum (Figure 3(b)) shows the $\alpha$ line strongly in emission, with an equivalent width of -22."
. The centre of the line corresponds to a velocity of 1654283 km consistent with SMC membership.," The centre of the line corresponds to a velocity of $\pm23$ km $^{-1}$, consistent with SMC membership."
 We have no photometry of objects in this field. but estimate 1—15 from Digitised Sky Survey images.," We have no photometry of objects in this field, but estimate $V \sim
15$ from Digitised Sky Survey images."
 This. together with the /fa emission and racial velocity indicates," This, together with the $H\alpha$ emission and radial velocity indicates"
It is also interesting to compare the Molfatian disk snapshots anc Newtonian ones.,It is also interesting to compare the Moffatian disk snapshots and Newtonian ones.
 Comparing Figures δ and 16.. one clearly sees how different these configurations are.," Comparing Figures \ref{fig8} and \ref{fig14}, one clearly sees how different these configurations are."
" Studies of alternative theories of gravity are usually made statically. instead of modeling ""live"" svstenis."," Studies of alternative theories of gravity are usually made statically, instead of modeling “live” systems."
 For example. in the study of spiral galaxies centrifugal equilibrium is usually considered.," For example, in the study of spiral galaxies centrifugal equilibrium is usually considered."
 The problem with this approach is that the secular evolution of these svstenis cannot be followed., The problem with this approach is that the secular evolution of these systems cannot be followed.
 In the present paper. in particular. we [follow the techniques used by Drandao&(2010a.b.c) and probe the DMs model. verifving if the Molfatian gravity is dynamically consistent to account for spiral galaxies.," In the present paper, in particular, we follow the techniques used by \citet{ca2009a,ca2009b,ca2009c} and probe the BM's model, verifying if the Moffatian gravity is dynamically consistent to account for spiral galaxies."
 A clvnamical study is a robust approach. since it considers the distribution functions and. Jeans equations that can maintain V-body svstems in equilibrium.," A dynamical study is a robust approach, since it considers the distribution functions and Jean's equations that can maintain $N$ -body systems in equilibrium."
 We performed the simulations with a modified and tested version of the Gadget-2 code., We performed the simulations with a modified and tested version of the Gadget-2 code.
 We replaced the Newtonian potential ancl acceleration bv (he expressions given by Equations (3)) and (4)) just in the respective code's instructions., We replaced the Newtonian potential and acceleration by the expressions given by Equations \ref{moffat_equation2}) ) and \ref{moffat_aceleration2}) ) just in the respective code's instructions.
 We have extensively tested (he code efficiency to caleulate potentials via the tree method. using a tvpical initial siapshot composed by a disk embedded in a dark matter halo.," We have extensively tested the code efficiency to calculate potentials via the tree method, using a typical initial snapshot composed by a disk embedded in a dark matter halo."
 From our simulations. it follows that the Mollatian potential cannot generate exponential disks in equilibrium. although the rotation curves are nearly flat throughout the whole simulated time.," From our simulations, it follows that the Moffatian potential cannot generate exponential disks in equilibrium, although the rotation curves are nearly flat throughout the whole simulated time."
 The final (stable) equilibrium configuration is very different [rom the exponential profile. as we have explained above.," The final (stable) equilibrium configuration is very different from the exponential profile, as we have explained above."
 So. if we consider Chat exponential profiles are reliable. as it seems (to be (he case. since (μον are consistent. wilh observations. ib ds hard to believe that DM model is better than the Newtonian model (barvonic disk," So, if we consider that exponential profiles are reliable, as it seems to be the case, since they are consistent with observations, it is hard to believe that BM model is better than the Newtonian model (baryonic disk"
be noted that neither of these galaxies have detectable levels of [OIIJ3727 emission which would indicate star formation and. furthermore. the galaxy. with lower surface brightness (CERSQ3.1071) is the bluer of the two galaxies. opposite to what might be expected if a recent. burst of star formation was a factor.,"be noted that neither of these galaxies have detectable levels of [OII]3727 emission which would indicate star formation and, furthermore, the galaxy with lower surface brightness (CFRS03.1077) is the bluer of the two galaxies, opposite to what might be expected if a recent burst of star formation was a factor."
 In summary. we have (wo independent methods of estimating the velocity dispersion of these two lens galaxies.," In summary, we have two independent methods of estimating the velocity dispersion of these two lens galaxies."
 The photometric method uses size and surlace brightness to estimate (he expected velocity dispersion assuming Chat these galaxies lie on either the local [fundamental plane or on an evolved fundamental plane., The photometric method uses size and surface brightness to estimate the expected velocity dispersion assuming that these galaxies lie on either the local fundamental plane or on an evolved fundamental plane.
 The lens method is independent of the galaxy. light and depends only on the lens geometry., The lens method is independent of the galaxy light and depends only on the lens geometry.
 The results of these two methods can be made to be consistent only if CFRSO3.L077 lies on the un-evolved fundamental plane whereas CFRS 14.1311 has evolved by ~1 magnitude at 2=0.307., The results of these two methods can be made to be consistent only if CFRS03.1077 lies on the un-evolved fundamental plane whereas CFRS 14.1311 has evolved by $\sim 1$ magnitude at $z=0.807$.
 In other words. the results suggest that these two hiehly luminous field elliptical galaxies maw have hacl very different evolutionary histories.," In other words, the results suggest that these two highly luminous field elliptical galaxies may have had very different evolutionary histories."
 Our spectroscopic observations confirm that (he arc surrounding (he z = 0.938 elliptical ealaxv CEIRS03.1077 is indeed a lensed image of a background galaxy., Our spectroscopic observations confirm that the arc surrounding the z = 0.938 elliptical galaxy CFRS03.1077 is indeed a lensed image of a background galaxy.
 The redshilt of this ealaxy dis z= 2.941., The redshift of this galaxy is z = 2.941.
 Standard lens models easily reproduce the observed are structure and also suggest Chat (wo faint objects observed near the lensing galaxy on the opposite side to the arc are lensed images., Standard lens models easily reproduce the observed arc structure and also suggest that two faint objects observed near the lensing galaxy on the opposite side to the arc are lensed images.
 Observations al other wavelengths should be obtained to determine the colors of these objects., Observations at other wavelengths should be obtained to determine the colors of these objects.
 If they. are the same as the arc. then (his would be further evidence that they are lensed images and can be conficlently used to constrain the lens geometry.," If they are the same as the arc, then this would be further evidence that they are lensed images and can be confidently used to constrain the lens geometry."
 Multiwavelength observations of CERS03.1077 could also be used to examine whether the infernal colors of the galaxy itself are normal or show strong variations indicative of recent star formation (cf Menanteau.Abraham&Ellis (2001)))., Multi-wavelength observations of CFRS03.1077 could also be used to examine whether the internal colors of the galaxy itself are normal or show strong variations indicative of recent star formation (cf \citet{Men01}) ).
 If the internal colors are nol homogeneous (his may help explain why the line of sight velocity dispersion. determined from the lens model is higher Caan expected from huidamental plane considerations assuming passive evolution since z~1., If the internal colors are not homogeneous this may help explain why the line of sight velocity dispersion determined from the lens model is higher than expected from fundamental plane considerations assuming passive evolution since $z \sim 1$.
 CFRSO3.1L077 demonstrates the potential offered by detailed study of Einstein. ring lenses., CFRS03.1077 demonstrates the potential offered by detailed study of Einstein ring lenses.
 With the advent of Sim class telescopes. especially those equipped with integral field spectrographs. the spectroscopic data reported here can now be very significantly improved.," With the advent of 8m class telescopes, especially those equipped with integral field spectrographs, the spectroscopic data reported here can now be very significantly improved."
 lmages at other wavelengths should be obtained to establish or identify additional lensed images in order to more tightly. constrain the lens mocel., Images at other wavelengths should be obtained to establish or identify additional lensed images in order to more tightly constrain the lens model.
" (Lissauer&Stevenson2007).. (Durisenetal.2007).. (Boley2009.andreferencestherein)... 2007).. iavan haveοιο,,TheArtvinowicz&Lubow1991:2008).. reduced"," \citep{lissauer07}, \citep{durisen07}. \citep[][and references therein]{boley09}. \citep[e.g.,][]{dm91, mathieu94,
 duchene07}. \citep[e.g.,][]{artymowicz94, ireland08}."
 (Deckwithotal.1990:Jensenet1900) (Chez," \citep{beckwith90, jensen96} \citep{ghez93, cieza09}."
 outer radius aud viscous timescale., outer radius and viscous timescale.
 These observational facts have ecnerally been interpreted as evidence that binaries tighter than ~100 AAU are mmuch less likely to support gas eiat. planet formation., These observational facts have generally been interpreted as evidence that binaries tighter than $\sim100$ AU are much less likely to support gas giant planet formation.
" However. follow-up nuaegnme surveys have ideutified some 50 plauct-lost stars that possess at least one stellar companion {οιοι,Pa-3009)."," However, follow-up imaging surveys have identified some 50 planet-host stars that possess at least one stellar companion \citep[e.g.,][]{patience02, chauvin06, raghavan06, eggenberger07,
 mugrauer09}."
 Tn particular. it it is worth noting that about of all known plaucts ii binary svstenis have a stellar colmpanion within less AAT. so that planet formation in such au cuviroument cannot be considered a rare occurrence.," In particular, it it is worth noting that about of all known planets in binary systems have a stellar companion within less AU, so that planet formation in such an environment cannot be considered a rare occurrence."
 Iu thisLetter. review several key statistical properties of PAIS aud field binary svsteuis that provide insight ou the planet formation process refsee:cd and 3)).," In this, I review several key statistical properties of PMS and field binary systems that provide insight on the planet formation process \\ref{sec:ci} and \ref{sec:end}) )."
 I then discuss the iuplications for the ain mechauisuis of plauet formation iu binary svstenis as a function of them projected separation refsecinplie))., I then discuss the implications for the main mechanisms of planet formation in binary systems as a function of their projected separation \\ref{sec:implic}) ).
 In this study. I only consider biuaries iu the AU. separation range. for which curreut PAIS inultipliitv surveys are reasonably complete.," In this study, I only consider binaries in the AU separation range, for which current PMS multiplicity surveys are reasonably complete."
 The tightest binary svsteni known to host a planot has a AAU separation., The tightest binary system known to host a planet has a AU separation.
 Stellar companions bevoud AAT are not expected to have πιο] influence ou planet formation., Stellar companions beyond AU are not expected to have much influence on planet formation.
 Iu order to draw a broad aud homogeneous view of the initial conditious for planet formation. I compiled ài. sample of 107 PAIS binaries for which deep (subjmiilliameter coutimmun observatious and/or neu to mid-intrared colors are available in the literature.," In order to draw a broad and homogeneous view of the initial conditions for planet formation, I compiled a sample of 107 PMS binaries for which deep (sub)millimeter continuum observations and/or near- to mid-infrared colors are available in the literature."
 The(subjmillameter data are taken from the work of Audrews&Williams(2005.20073: for almost all tarects. a lo sensitivity of nunJy. or better at 850400 and/or l.9uuu is achieved.," The (sub)millimeter data are taken from the work of \citet{andrews05,andrews07}; for almost all targets, a $\sigma$ sensitivity of mJy or better at $\mu$ m and/or 1.3mm is achieved."
 The imedian projectedseparation in this sample is AAT., The median projectedseparation in this sample is AU.
 I also defined a comparison suuple of 222 PAIS stars for which no companion, I also defined a comparison sample of 222 PMS stars for which no companion
drawn Irom the updated version of Harris (1996) and from Mackey van den Dergh (2005).,drawn from the updated version of Harris (1996) and from Mackey van den Bergh (2005).
 The data in both of these tables are. of course. most incomplete for intrinsically [aint elobulars.," The data in both of these tables are, of course, most incomplete for intrinsically faint globulars."
" The Andromeda galaxy is seen to contain 16 clusters brighter than M, = -10.0. compared to only 2 such clusters in the Milkv Way System."," The Andromeda galaxy is seen to contain 16 clusters brighter than $M_{v}$ = -10.0, compared to only 2 such clusters in the Milky Way System."
" The corresponding figures for clusters with -9.99 <M, -9.00 and -8.99 <Af, -8.00 are 35 and 10. ancl 88 ancl 26. respectively."," The corresponding figures for clusters with -9.99 $< M_{v} <$ -9.00 and -8.99 $< M_{v} <$ -8.00 are 35 and 10, and 88 and 26, respectively."
 If the luminosity distributions of the elobular clusters are similar in (hese (wo galaxies. then these results indicate (hat the M31 cluster svstem contains about [our times as many elobular clusters as does its Galactic counterpart.," If the luminosity distributions of the globular clusters are similar in these two galaxies, then these results indicate that the M31 cluster system contains about four times as many globular clusters as does its Galactic counterpart."
 A number of investigations of the dynamical evolution of elobular clusters (Spitzer Thuan 1972. Lightman Shapiro 1978. Murphy. Cohn Int 1990. Aarseth Iegeie 1993) have shown that the hall-light radii Hj of such clusters evolve very little over periods as long as ten cluster relaxation times.," A number of investigations of the dynamical evolution of globular clusters (Spitzer Thuan 1972, Lightman Shapiro 1978, Murphy, Cohn Hut 1990, Aarseth Heggie 1998) have shown that the half-light radii $R_{h}$ of such clusters evolve very little over periods as long as ten cluster relaxation times."
 The cluster hall-light radius is therefore a useful parameter for (he study of possible intrinsic differences between (he globular clusters in different. galaxies., The cluster half-light radius is therefore a useful parameter for the study of possible intrinsic differences between the globular clusters in different galaxies.
 In the discussion given below use will be made of the median of the half-leht radii of various data sets., In the discussion given below use will be made of the median of the half-light radii of various data sets.
 Beers et al. (, Beers et al. (
1990) have shown that this parameter is of comparable accuracy (ο more complicated metrics for populations with n > 100.,1990) have shown that this parameter is of comparable accuracy to more complicated metrics for populations with n $>$ 100.
 The possibility that the radii of elobular clusters might be a useful parameter for the estimation of distances to clusters was first hinted at by Shaplev and Sawyer (1927)., The possibility that the radii of globular clusters might be a useful parameter for the estimation of distances to clusters was first hinted at by Shapley and Sawyer (1927).
 The fact that the radii of globular clusters are in fact (at least within the main body of the Milkv. Wav) independent of their huminosities was [ist established by van den Bereh Alorbev (1984)., The fact that the radii of globular clusters are in fact (at least within the main body of the Milky Way) independent of their luminosities was first established by van den Bergh Morbey (1984).
 This conclusion is confirmed by the data on the globular clusters in the Andromeda galaxy and in the Milkv. Way System that are discussed below., This conclusion is confirmed by the data on the globular clusters in the Andromeda galaxy and in the Milky Way System that are discussed below.
 The result that the hall-lHight radii of elobular clusters associated with earlv-tvpe galaxies are independent of elobular cluster, The result that the half-light radii of globular clusters associated with early-type galaxies are independent of globular cluster
the Llp fraction. by mass. of the total eas.,"the $_2$ fraction, by mass, of the total gas."
 Phe first peak. between 2.0 and 2.1 kpe. is interpreted as newly formed. LL» that has been formed in the vicinity of the shock at that site in the SPA.," The first peak, between 2.0 and 2.1 kpc, is interpreted as newly formed $_2$ that has been formed in the vicinity of the shock at that site in the SPA."
 The other two high-LI» areas bevond the shock. near 2.25 and 2.6 kpe. would be (giant) molecular. clouds that were formed upstream by the same shock mechanism. at higher longitude. and which are moving through the SPA.," The other two $_2$ areas beyond the shock, near 2.25 and 2.6 kpc, would be (giant) molecular clouds that were formed upstream by the same shock mechanism, at higher longitude, and which are moving through the SPA."
 In the absence of feedback in our simulation. these clouds may be larger than expected. and longer-Iived.," In the absence of feedback in our simulation, these clouds may be larger than expected, and longer-lived."
 The outer edge ofthe SPA is less easy to define than the inner edge. given that there is no sharp outer boundary. but across much of the SPA. the density begins to decrease toward interarm values approximately 0.7 to 0.8 kpe behind the inner boundary.," The outer edge of the SPA is less easy to define than the inner edge, given that there is no sharp outer boundary, but across much of the SPA, the density begins to decrease toward interarm values approximately 0.7 to 0.8 kpc behind the inner boundary."
 The effect. of these density and temperature variations on the observed profile can be informed by the lower right panel of figure 10.. which shows the intensity contribution d/ per unit of column density as a function of distance.," The effect of these density and temperature variations on the observed profile can be informed by the lower right panel of figure \ref{fighisa}, which shows the intensity contribution $dI$ per unit of column density as a function of distance."
 LUISA (negative dL) appears to be produced behind the shock. once the temperature has dropped to below LOO Ix. and where the density is above LO7? g 7.," HISA (negative $dI$ ) appears to be produced behind the shock, once the temperature has dropped to below 100 K, and where the density is above $10^{-23}$ g $^{-3}$."
 At the front of the shock rere is a narrow region where emission is produced rather wn absorption. because the gas is still quite warm. vet the density is already increasing rapidly.," At the front of the shock there is a narrow region where emission is produced rather than absorption, because the gas is still quite warm, yet the density is already increasing rapidly."
 Such conditions are avourable for the onset of Le production. which depends μαronely on density as well as weakly on temperature (Berginal. 2004).," Such conditions are favourable for the onset of $_2$ production, which depends strongly on density as well as weakly on temperature \citep{Ber04}."
. Still. this shows that within the region where 10 gas velocity is stronely influcneed by the spiral acm shock. both emission and absorption fromLLL atoms are xossible.," Still, this shows that within the region where the gas velocity is strongly influenced by the spiral arm shock, both emission and absorption from atoms are possible."
 Yet we can constrain the occurrence of LISA to cold. dense regions which also contain molecular gas.," Yet we can constrain the occurrence of HISA to cold, dense regions which also contain molecular gas."
 The emperature appears to remain at values around. 100 Ix out ο a distance of ~2.65 kpe. after which it begins to rise anc he molecular fraction drops sharply.," The temperature appears to remain at values around 100 K out to a distance of $\sim 2.65$ kpc, after which it begins to rise and the molecular fraction drops sharply."
 Thus the cold gas is mostly confined to the spiral arms. ancl the eas is warmed in he interarm regions.," Thus the cold gas is mostly confined to the spiral arms, and the gas is warmed in the interarm regions."
 Dobbsetal.(2008) found that roughly 62 per cent of the was in the cold phase. where 7<150 Ix was their criterion for cold gas.," \citet{DobbsGloverClarkKlessen2008} found that roughly 62 per cent of the was in the cold phase, where $T \le 150$ K was their criterion for cold gas."
 The shock region itself is much smaller than the ~100 , The shock region itself is much smaller than the $\sim 100$ 
to Nee tuplics that the most luuinous objects are collapsing.,to $_{CO}$ implies that the most luminous objects are collapsing.
 Braud Wouterloot (1995) examined the variation of the couversion factor using aand oobservatious for a saaple of far outer Galaxy clouds., Brand Wouterloot (1995) examined the variation of the conversion factor using and observations for a sample of far outer Galaxy clouds.
 Accounting for a radial gradient in aabundance. they concluded that the conversion factor is similar to that found in the inner Galaxy.," Accounting for a radial gradient in abundance, they concluded that the conversion factor is similar to that found in the inner Galaxy."
 Other aand sstudies have found no significant variation of Ney for outer Galaxy clouds (Carpenter. Suell. Schloerb L990).," Other and studies have found no significant variation of $X_{CO}$ for outer Galaxy clouds (Carpenter, Snell, Schloerb 1990)."
 The conventional assumption of a constant conversion factor is that clouds are self gravitationally bound with ας21., The conventional assumption of a constant conversion factor is that clouds are self gravitationally bound with $\alpha_G \approx 1$.
 However. as discussed in 52.1. if clouds are internally overpressured with respect to self eravity (ας22 1). then the appropriate value of Vee is snaller than the standard. constant value.," However, as discussed in $\S$ 2.4, if clouds are internally overpressured with respect to self gravity $\alpha_G >>1$ ), then the appropriate value of $X_{CO}$ is smaller than the standard, constant value."
 Therefore. * using the staudard value. the derived masses are upper lanits aud the derived values for ag; are lower nuits.," Therefore, by using the standard value, the derived masses are upper limits and the derived values for $\alpha_G$ are lower limits."
 The huge values of the eravitational parameter refiect the changing dynamical state of molecular reeions with different mass., The large values of the gravitational parameter reflect the changing dynamical state of molecular regions with different mass.
 Only the most Iuninous objects identified in the Survey lave sufficient mass to )e bound by self eravity (oc; ~1)., Only the most luminous objects identified in the Survey have sufficient mass to be bound by self gravity $\alpha_G\sim$ 1).
 Regions with lower CO luminosities and mass are internally overpressured with respect to self eravity., Regions with lower CO luminosities and mass are internally overpressured with respect to self gravity.
 This state is independent of whether the object is au isolated. cloud or part of a larger cloud complex., This state is independent of whether the object is an isolated cloud or part of a larger cloud complex.
 A similar conclusion has been obtained for a sample of high latitude clouds and or several clouds in the solar uciglborhood (Maguaui. Blitz. Muudy 1985: keto Myers 1986: Bertoldi Melsoe 1992: Falearouc. Puget. Perault 1992: Dobashi 1996: Yonekura 1997: Iuviunura 1998).," A similar conclusion has been obtained for a sample of high latitude clouds and for several clouds in the solar neighborhood (Magnani, Blitz, Mundy 1985; Keto Myers 1986; Bertoldi McKee 1992; Falgarone, Puget, Perault 1992; Dobashi 1996; Yonekura 1997; Kawamura 1998)."
 The results preseuted here provide statistical confirmation of these earlier studies over a larger range of cloud and chump masses., The results presented here provide statistical confirmation of these earlier studies over a larger range of cloud and clump masses.
 Iu order to gauge the results of the previous section with amore reliable tracer of molecular hydrogen mass. we have analyzed aaud ddata of targeted molecular cloud regions which lie within the Survey field (Hover. Carpeuter. Ladd 1996: Deane 2000).," In order to gauge the results of the previous section with a more reliable tracer of molecular hydrogen mass, we have analyzed and data of targeted molecular cloud regions which lie within the Survey field (Heyer, Carpenter, Ladd 1996; Deane 2000)."
" The targeted ficlds include the giant molecular clouds Cep OB3. 5110, NGC 7538, and W3."," The targeted fields include the giant molecular clouds Cep OB3, S140, NGC 7538, and W3."
 The ddata were decomposed into discrete objects with the same algorithin as the Survey cube., The data were decomposed into discrete objects with the same algorithm as the Survey cube.
 The Bnuteerated intensity is sununed within the boundaries identified from the ddata and a ass {τε is derived assuuniüug local thermodvuamic equilibrium. a kinetic temperature of 10 Kk. a to aabundauce of G aud a 1.36 correction for the abundance of Ποπ (Dickman 1978).," The integrated intensity is summed within the boundaries identified from the data and a mass, $M_{LTE}$, is derived assuming local thermodynamic equilibrium, a kinetic temperature of 10 K, a to abundance of $^{-6}$ and a 1.36 correction for the abundance of Helium (Dickman 1978)."
" The distauces to each cloud ave: 730 pe (Cop OB3). 910 pe (S110). 2.35 kpe (V3). and 3.5 kpe (NGC 7538),"," The distances to each cloud are: 730 pc (Cep OB3), 910 pc (S140), 2.35 kpc (W3), and 3.5 kpc (NGC 7538)."
 A virial mass or each object is derived from the tabulated size and vvolocitv dispersion., A virial mass for each object is derived from the tabulated size and velocity dispersion.
 The variation of the gravitationalparameter. now derived with unieasurenments of molecular cola density aud Appr. is shown in Figure 5 for the four targeted eiaut uolecular clouds.," The variation of the gravitationalparameter, now derived with measurements of molecular column density and $M_{LTE}$, is shown in Figure \ref{alphaG13co} for the four targeted giant molecular clouds."
" The evaluation of ας is subject to the same selection effects as the Survey clouds such hat there is the same fictional depeudeuce of az""miu with Aare.", The evaluation of $\alpha_G$ is subject to the same selection effects as the Survey clouds such that there is the same functional dependence of $\alpha_G^{min}$ with $M_{LTE}$.
 oc; decreases with increasing Iuniuositv and mass., $\alpha_G$ decreases with increasing luminosity and mass.
 Bisector fits of the data to the expression where AL. is as defiued iu equation 23.1. are sumunarizecd in Table 2.," Bisector fits of the data to the expression where $_\circ$ is as defined in equation 3.4, are summarized in Table 2."
 Values of e range between 0.51 and 0.58., Values of $\epsilon$ range between 0.51 and 0.58.
 For objects with masses exeater than 105 AZ... the derived values of àc; are reasonably consistent with," For objects with masses greater than $^4$ , the derived values of $\alpha_G$ are reasonably consistent with"
we therefore exclude it from our catalog.,we therefore exclude it from our catalog.
 The values for the sharpness and crowding cuts for ΕΤΟΣ were the standard used by the ANGST. program., The values for the sharpness and crowding cuts for WFPC2 were the standard used by the ANGST program.
 The cuts for ΝΕΟΣ were chosen by looking at CMDs resulting from different possibilities aud choosing values that produced clean CMD features and a low umber of contaminants. falling outside of auv known feature.," The cuts for WFC3 were chosen by looking at CMDs resulting from different possibilities and choosing values that produced clean CMD features and a low number of contaminants, falling outside of any known feature."
 We note that our tests of different cuts showed that the choices of cuts have little inuipact on our final SEIT micasireients outside of the most recent time bin (<10 Myr). most likely due to the significant clustering of the very vouugest stars.," We note that our tests of different cuts showed that the choices of cuts have little impact on our final SFH measurements outside of the most recent time bin $<$ 10 Myr), most likely due to the significant clustering of the very youngest stars."
 Our final optical star catalogs contained 16806. 28088. 138350. and stars for the outer (WEDPC2 N). bridge (WEDPC2 NW). and ceutral (UVIS) fields. respectively.," Our final optical star catalogs contained 16806, 28088, 138350, and stars for the outer (WFPC2 N), bridge (WFPC2 NW), and central (UVIS) fields, respectively."
 Our final UV catalog for the central field contained T7656 stars. and our final IR catalog for the ceutral feld contained 19807 stars.," Our final UV catalog for the central field contained 77656 stars, and our final IR catalog for the central field contained 19807 stars."
 In the eud. our analysis of the optical (F13sW aud. ESLIW) plotometiy frou UVIS vielded results that were consistent with those of the UV (F336W aud F138W) photometry from UVIS.," In the end, our analysis of the optical (F438W and F814W) photometry from UVIS yielded results that were consistent with those of the UV (F336W and F438W) photometry from UVIS."
 Since no additional insight was eained from the optical photometry. we will uot discuss it further but imclude the ΕΙ aud limiting maenitucdes for reference.," Since no additional insight was gained from the optical photometry, we will not discuss it further but include the SFH and limiting magnitudes for reference."
 Any statements made about the UW photometry from UVIS was confirmed with the optical photometry frou UVIS., Any statements made about the UV photometry from UVIS was confirmed with the optical photometry from UVIS.
 The final €CMDs are shown in Figure 2.., The final CMDs are shown in Figure \ref{cmds}.
 While all the WEDPC?2 tuagineg of the outer portions of the disk provided high-quality photometry goie much fainter than the tip of the red elaut branch (and. even fainter than the red clump in the deep field). the ceutral ΝΕΟΣ fields were not as consistent.," While all the WFPC2 imaging of the outer portions of the disk provided high-quality photometry going much fainter than the tip of the red giant branch (and even fainter than the red clump in the deep field), the central WFC3 fields were not as consistent."
 The UVIS images proved to lave similar completcuess limits tleoughout the field. but the IR nuages were clearly overcrowded near the galaxy ceuter.," The UVIS images proved to have similar completeness limits throughout the field, but the IR images were clearly over-crowded near the galaxy center."
 This strong differeutial crowding in the IR data mace aualvsis of the full IR feld dubious., This strong differential crowding in the IR data made analysis of the full IR field dubious.
 We therefore separated the iuuennost portion of the central field into two spatial components. following the isophotes of the galaxy.," We therefore separated the innermost portion of the central field into two spatial components, following the isophotes of the galaxy."
 Our division cllipse and the ealaxy isophotes iu the FLIOW nage are shown iu Figure 3.., Our division ellipse and the galaxy isophotes in the F110W image are shown in Figure \ref{ellipses}.
 We measured the star formation rate and metallicity as a function of stellar age using the software package AMIATCT to fit the observed CAIDs., We measured the star formation rate and metallicity as a function of stellar age using the software package MATCH to fit the observed CMDs.
 We adopted magnitude cuts set to the limits provided iu Table 1.. and then fitted the stellar evolution models of andmainBodyCitationEud1e¢irardi2002. convolved with our photometric error aud703] completcuess statistics and populated with a initial mass function (IMIF).," We adopted magnitude cuts set to the limits provided in Table \ref{table}, and then fitted the stellar evolution models of and, convolved with our photometric error and completeness statistics and populated with a initial mass function (IMF)."
 The choices of software aud models used for the ANGST project are discussed in detail iu(2009b)., The choices of software and models used for the ANGST project are discussed in detail in.
. Briefly. comparisons of results when fitting CMDs using different models aud fitting software have shown that the results are consistent within the estimated uucertaimties: furthermore. our choice of models provides the largest range of ages and muctallicities publicly available.," Briefly, comparisons of results when fitting CMDs using different models and fitting software have shown that the results are consistent within the estimated uncertainties; furthermore, our choice of models provides the largest range of ages and metallicities publicly available."
 We first ft the data assuming a single foreground reddening Ay-=0.07 and distance im Afyg=27.11. which were obtained from the ANGST survev2009).. aud the Galactic dust maps.," We first fit the data assuming a single foreground reddening $A_V$ =0.07 and distance $m-M_0$ =27.41, which were obtained from the ANGST survey, and the Galactic dust maps."
 The best fit provides the relative contribution of stars of cach age aud ietallicity in each field., The best fit provides the relative contribution of stars of each age and metallicity in each field.
" For the shallowest data. the IR and optical data for the iunerinost region. we limited the nmuber of free parameters by iuposiug an ""juereasne inctallicity” coustraiut on the fit. whereby the metallicity of the population was not allowed to decrease with time (within the measured errors)."," For the shallowest data, the IR and optical data for the innermost region, we limited the number of free parameters by imposing an “increasing metallicity” constraint on the fit, whereby the metallicity of the population was not allowed to decrease with time (within the measured errors)."
 These fits are turued iuto SEIIS of star formation rate ancl metallicitv as à function of time., These fits are turned into SFHs of star formation rate and metallicity as a function of time.
 To assess the uncertainties of our fits. we then rau Monte Carlo fitting tests.," To assess the uncertainties of our fits, we then ran Monte Carlo fitting tests."
 These tests assess 2 types of errors: random errors due to Poisson sampling auc errors in photometry. aud svstematic errors due to deficiencies in the stellar evolution models as well as uv offset iu distauce. reddening. and/or zero-poiuts.," These tests assess 2 types of errors: random errors due to Poisson sampling and errors in photometry, and systematic errors due to deficiencies in the stellar evolution models as well as any offset in distance, reddening, and/or zero-points."
 The Poisson errors are accounted for bv resampling the couvolved vest-fit inodel LOO times., The Poisson errors are accounted for by resampling the convolved best-fit model 100 times.
 Then. when fitting each of hese realizations. the systematic errors are accounted for by introducing stall raucdom shifts iu the bolometric uaenitudes and effective. temperatures of the models.," Then, when fitting each of these realizations, the systematic errors are accounted for by introducing small random shifts in the bolometric magnitudes and effective temperatures of the models."
 These shifts are introduced at the level of the differences οποσα models in the literature. aud therefore serve as a proxy of the effects of our choice of stellar evolution uodels.," These shifts are introduced at the level of the differences between models in the literature, and therefore serve as a proxy of the effects of our choice of stellar evolution models."
 Our final uncertainties were calculated as the confidence iutervals of the results of all of our Monte Carlo test fits., Our final uncertainties were calculated as the confidence intervals of the results of all of our Monte Carlo test fits.
 These total uncertainties are used as the error bars in all subsequent plots and analysis., These total uncertainties are used as the error bars in all subsequent plots and analysis.
 We can also assess the deeree of systematic errors using plots of residuals between the observed CMD aud the model fit., We can also assess the degree of systematic errors using plots of residuals between the observed CMD and the model fit.
 An example of the residuals of our fits is shown in Figure L., An example of the residuals of our fits is shown in Figure \ref{residuals}.
 Our fits assuue the Calactic foreground extinction over the field., Our fits assume the Galactic foreground extinction over the field.
 Therefore. the low-level residuals over the entire CMD for these fits suggests that our photometry is dominated bv stars in frout of any significant internal dust laver iu the NCC 1211 disk. with a modest amount of internal extinction (cf. 2?))," Therefore, the low-level residuals over the entire CMD for these fits suggests that our photometry is dominated by stars in front of any significant internal dust layer in the NGC 4214 disk, with a modest amount of internal extinction (cf. \ref{bright}) )"
 causing the slieht excess in the data redward of the model main-sequence., causing the slight excess in the data redward of the model main-sequence.
 Furthermore. the lack of aux strong features in the fit residuals suggests that the distribution of stellar luinositics is well-fitted assuimiug a Salpeter IME. once the effects of the age distribution are modeled.," Furthermore, the lack of any strong features in the fit residuals suggests that the distribution of stellar luminosities is well-fitted assuming a Salpeter IMF, once the effects of the age distribution are modeled."
 These fits show that the main sequence lIumunositv profile iterpreted by as sueecsting variations in the IME. is also consistent with a standard IME with a varvine star formation lustory: thus IMF variations and SEIT are somewhat degenerate in such a complex region.," These fits show that the main sequence luminosity profile interpreted by as suggesting variations in the IMF, is also consistent with a standard IMF with a varying star formation history; thus IMF variations and SFH are somewhat degenerate in such a complex region."
 Our uncertainties conie froni iueasurenients assunüneg a constant IME., Our uncertainties come from measurements assuming a constant IMF.
" ""Therefore. our uncertainties would increase if variations in the IME were allowed."," Therefore, our uncertainties would increase if variations in the IMF were allowed."
 We used a very fine eri of model stellar isochrones (0.05 dex iu age aud 0.1 dex in |Fe/II]) to provide the best possible fit to our data., We used a very fine grid of model stellar isochrones (0.05 dex in age and 0.1 dex in [Fe/H]) to provide the best possible fit to our data.
 Because the erid is very fine. pushing star formation between adjaceut time bius has little effect on the quality of the resulting fit.," Because the grid is very fine, pushing star formation between adjacent time bins has little effect on the quality of the resulting fit."
 The degree to which this degeneracy is true depends on the quality of the data beiug fitted. which is characterized by our Moute Carlo tests.," The degree to which this degeneracy is true depends on the quality of the data being fitted, which is characterized by our Monte Carlo tests."
 The uncertainties m the cumulative age distribution cannot be improved bv further binning of the data in time., The uncertainties in the cumulative age distribution cannot be improved by further binning of the data in time.
 Therefore. we plot the full time resolution of," Therefore, we plot the full time resolution of"
that the hosts al z«1 are not representative of the general galaxy population(Levesque el al.,that the hosts at $z<1$ are not representative of the general galaxy population(Levesque et al.
 20090)., 2009a).
 Thus. the properties of these low-redshift GRB hosts presented in this paper could not be reproduced by anv monolithic process.," Thus, the properties of these low-redshift GRB hosts presented in this paper could not be reproduced by any monolithic process."
 At least. some low-redshift galaxies maa undergo multiple star-Iorming processes during their whole liletimes.," At least, some low-redshift galaxies may undergo multiple star-forming processes during their whole lifetimes."
 GRB production can be accompanied with any single starburst event., GRB production can be accompanied with any single starburst event.
 From the analvsis in this paper. we see that (he absorption of GRB X-ray. and optical emissions is relatively strong.," From the analysis in this paper, we see that the absorption of GRB X-ray and optical emissions is relatively strong."
" The strong intrinsic attenuation of GRB host galaxies may produce some clark bursts. defined by the index »,,<0.5. where ο, is the flux density ratio between optical ancl X-ray bands(Jakobsson et al."," The strong intrinsic attenuation of GRB host galaxies may produce some dark bursts, defined by the index $\beta_{ox}<0.5$ , where $\beta_{ox}$ is the flux density ratio between optical and X-ray bands(Jakobsson et al."
 2004)., 2004).
 Rol et al. (, Rol et al. (
2005) proposed several ex(netion origins from their preliminary results.,2005) proposed several extinction origins from their preliminary results.
 From our calculations. we see that the heavy altenuation may occur due to the following three possibilities: (1) the local environment of the 10st is metal-enriched. metallicity is hieher. and/or. the host galaxy in (he massive dark halo arger than LO’AL. may have strong absorption.," From our calculations, we see that the heavy attenuation may occur due to the following three possibilities: (1) the local environment of the host is metal-enriched, metallicity is higher, and/or, the host galaxy in the massive dark halo larger than $10^{12}M_\odot$ may have strong absorption."
" For example. at redshilt 2.5. Z=1.0Z.. ido mass ρω=5.0x LOPAL.. after the galactic evolving time 1.0x105yr. we have dust extinction 24,=1.0 and the corresponding X-ray absorption Ny,=4.7xLO?en7: (3) the dust ancl metals surrounding the GRB in the host galaxy are distributed in an inhomogeneous wav: there could be heavy absorption through the line of sight. but in other directions the absorption is slight."," For example, at redshift 2.5, $Z=1.0Z_\odot$, halo mass $M_{halo}=5.0\times 10^{12}M_{\odot}$ , after the galactic evolving time $1.0\times 10^8~yr$, we have dust extinction $A_v=1.0$ and the corresponding X-ray absorption $N_{H,x}=4.7\times 10^{22}cm^{-2}$; (2) the dust and metals surrounding the GRB in the host galaxy are distributed in an inhomogeneous way; there could be heavy absorption through the line of sight, but in other directions the absorption is slight."
" Also. in our model. we assume that the sl, and Nyy, are neasured locally and do not change significantly if the dust and gas extend out to a few tens to hundreds of pe from the burst(Perna Lazzati 2002. D'Elia et al."," Also, in our model, we assume that the $A_v$ and $N_{H,x}$ are measured locally and do not change significantly if the dust and gas extend out to a few tens to hundreds of pc from the burst(Perna Lazzati 2002, D'Elia et al."
 2009): however. suppose the observed optical extinction is due to the grain absorption lar bevond (his local region ol GRD. the τμ correlation obtained by Schady et al. (," 2009); however, suppose the observed optical extinction is due to the grain absorption far beyond this local region of GRB, the $A_v$ $N_{H,x}$ correlation obtained by Schady et al. ("
2007. 2010) may be invalid and our calculations are stronglv biased: (3) as mentioned in Section 2.2.2. the dust produced by ihe AGD population at high redshift should be taken into account.,"2007, 2010) may be invalid and our calculations are strongly biased; (3) as mentioned in Section 2.2.2, the dust produced by the AGB population at high redshift should be taken into account."
 In order to further understancl the metal production of GRB environment. we roughly re-estimate the metallicity of GRD hosts under our framework.," In order to further understand the metal production of GRB environment, we roughly re-estimate the metallicity of GRB hosts under our framework."
" The mass of metal ως. where fis the ratio of massive stars to all stars. and. fy, is the ratio of mental converted. from the dust."," The mass of metal $M_{metal}=SFR\cdot f\cdot f_{dep}\cdot M_{dust}/M_{star}$ , where f is the ratio of massive stars to all stars, and $f_{dep}$ is the ratio of mental converted from the dust."
 /=0.47 is the case for the stars with the mass larger than 2.M. by our adopted IMF. fae)—1.0 means that all the dust can be transferred to metals.," $f=0.47$ is the case for the stars with the mass larger than $2M_\odot$ by our adopted IMF, $f_{dep}=1.0$ means that all the dust can be transferred to metals."
" Metallicity is defined by Z=M,,;,,/Mu.", Metallicity is defined by $Z=M_{metal}/M_{gas}$.
 From the SER calculated by Granato et al. (, From the SFR calculated by Granato et al. (
2004) and. Mao et al. (,2004) and Mao et al. (
2007). as an example. al redshift 6. we obtain the metallicity as Zc2555x10ΑμMass).,"2007), as an example, at redshift 6, we obtain the metallicity as $Z\sim 2.75\times 10^{-2}(M_{dust}/M_{star})$."
 HE we take a supernova with the dust production of 10.ᾗ solar mass(Pozzo et al., If we take a supernova with the dust production of $10^{-3}$ solar mass(Pozzo et al.
 2004). we obtain the upper limit of metallicity Z~107Z.. which is lower (han the measurement(Z> 0.02Z.) of GRD 050904(Canmpana et al.," 2004), we obtain the upper limit of metallicity $Z\sim 10^{-3}Z_{\odot}$, which is lower than the $Z>0.02Z_\odot$ ) of GRB 050904(Campana et al."
 2007)., 2007).
 IE we take the dust mass 0.08-0.3 solar mass per primordial massive supernova(Todini Ferrara. 2001). we have the result which isconsistent with the observation.," If we take the dust mass 0.08-0.3 solar mass per primordial massive supernova(Todini Ferrara 2001), we have the result which isconsistent with the observation."
The estimation values of Population L/II metallicity are lower (han (he observational values at high redshift. meaning,"The estimation values of Population I/II metallicity are lower than the observational values at high redshift, meaning"
no evidence for a substantial increase in optical obscuration in the galaxy population between 250 and z=0.5.,no evidence for a substantial increase in optical obscuration in the galaxy population between z=0 and z=0.5.
 We have. though. detected. strong subnim emission. from one object. in addition to the one strong detection. ancl one weaker detection from Paper 1.," We have, though, detected strong submm emission from one object, in addition to the one strong detection and one weaker detection from Paper 1."
 LIST imaging of the two bright submin sources suggests they are quiescent. edee on disk galaxies. in contrast to the disturbed. ancl interacting SAIGs. found in. blank Ποιά submnm surveys. and local ULIRGs.," HST imaging of the two bright submm sources suggests they are quiescent, edge on disk galaxies, in contrast to the disturbed and interacting SMGs, found in blank field submm surveys, and local ULIRGs."
 We examine the dust. properties of these sources. ancl conclude that hey are ultra-Iuminous submini sources. with large dust masses. but with cold cust SEDs. T-—20k. We also examine the luminosity distribution of the SNIa host galaxies as à €ass. and compare it to the local SLUGS submim luminosity function.," We examine the dust properties of these sources, and conclude that they are ultra-luminous submm sources, with large dust masses, but with cold dust SEDs, $\sim$ 20K. We also examine the luminosity distribution of the SN1a host galaxies as a class, and compare it to the local SLUGS submm luminosity function."
 We conclude that the z=0.5 luminosity clistribution seems bimodal. with an absence of sources within a factor of a few of our bright objects.," We conclude that the z=0.5 luminosity distribution seems bimodal, with an absence of sources within a factor of a few of our bright objects."
 Several scenarios could explain this effect. including some forms of luminosity function evolution. the absence of sources such as these from local submim surveys. or problems with the classification of 1 of SNe in the twpe la SN surveys.," Several scenarios could explain this effect, including some forms of luminosity function evolution, the absence of sources such as these from local submm surveys, or problems with the classification of $\sim$ of SNe in the type 1a SN surveys."
 We propose various wavs in which these hyptjeses could be tested by future observations., We propose various ways in which these hyptheses could be tested by future observations.
 TheJames Clerk Maxwell Telescope is operated. by Vhe Joint Astronomy Centre on behalf. of the Particle hesies ancl Astronomy Research Council of the United ]xingdonm.D the Netherlands Oreanisation5 [or Scientific Lescarch. and the National Research Council of Canada.," The James Clerk Maxwell Telescope is operated by The Joint Astronomy Centre on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada."
 DLC acknowledges funcing from PPARC., DLC acknowledges funding from PPARC.
" JA gratefully acknowledges the support from the Science and Technology Foundation (ECT. Portugal) through the research erant ""OCTPI-ENU-43805-2001."," JA gratefully acknowledges the support from the Science and Technology Foundation (FCT, Portugal) through the research grant POCTI-FNU-43805-2001."
 The authors would like to thank he stalf at the ΤΙ for their usual excellent: support work. and Ro Wotak for useful discussions.," The authors would like to thank the staff at the JCMT for their usual excellent support work, and R. Kotak for useful discussions."
" “Phe research deseribed in this paper was carried out. in part. by the Jet ""propulsion. Laboratory. California Institute of Technology. and was sponsored by theNational Xeronauties and Space Administration."," The research described in this paper was carried out, in part, by the Jet Propulsion Laboratory, California Institute of Technology, and was sponsored by the National Aeronautics and Space Administration."
 This publication is also based in part on observations made with the NASA/ESA IIubble Space Telescope. obtained. from the Data Archive at the Space Telescope Science Institute. which is operated by the Association of Universities for neearch in Astronomy. Lne.. under NASA contract NASA5-26555.," This publication is also based in part on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NASÊ5-26555."
are grateful to D. Marrone. S. Horustein aud F. Baganoll for provicing us with their data.,"are grateful to D. Marrone, S. Hornstein and F. Baganoff for providing us with their data."
"the bias measurement to estimate the effectiveness of the method for example galaxy populations, i-dropouts at z~6 in ACS-sized fields and F600LP-dropouts at z~7 in WFC3-sized fields.","the bias measurement to estimate the effectiveness of the method for example galaxy populations, $i$ -dropouts at $z\sim6$ in ACS-sized fields and $\FSIXLP$ -dropouts at $z\sim7$ in WFC3-sized fields."
 Our method can be for other dropout filter strategies., Our method can be adapted for other dropout filter strategies.
" At z~6 (z~ 7) the adaptedmethod should provide a measurement of galaxy bias with SNR=3 for luminosities near Muv7M, when Nsgaas>40 (Ngeus2, 130) fields are used.", At $z\sim6$ $z\sim7$ ) the method should provide a measurement of galaxy bias with $SNR\gtrsim3$ for luminosities near $\MUV\approx\Mstar$ when $\Nfields\ge40$ $\Nfields\gtrsim130$ ) fields are used.
 The uncertainty on the galaxy bias improves with an increasing number of independent fields as £z1//Ngeias., The uncertainty on the galaxy bias improves with an increasing number of independent fields as $\approx1/\sqrt{\Nfields}$.
" These requirements on the number of fields are concomitant with the expected number of pointings available from extensions to ongoing pure parallel HST programs (Trenti2008;Yan2008), and the potential for using this method to measure high-redshift galaxy bias is therefore promising."," These requirements on the number of fields are concomitant with the expected number of pointings available from extensions to ongoing pure parallel HST programs \citep[][]{trenti2008b,yan2008a}, and the potential for using this method to measure high-redshift galaxy bias is therefore promising."
" We also show that using correlated, nearby fields to perform the presented measurement typically leads to an underestimate of the bias owing to a correlation-induced decrease in the dispersion of galaxy number counts."," We also show that using correlated, nearby fields to perform the presented measurement typically leads to an underestimate of the bias owing to a correlation-induced decrease in the dispersion of galaxy number counts."
 The measured bias can be associated with a characteristic mass of DM halos with similar clustering., The measured bias can be associated with a characteristic mass of DM halos with similar clustering.
 We show that the typical uncertainty in the bias measured using the presented method corresponds to an uncertainty in the inferred DM halo mass of ~ a few percent in logM for i-dropouts., We show that the typical uncertainty in the bias measured using the presented method corresponds to an uncertainty in the inferred DM halo mass of $\sim$ a few percent in $\log M$ for $i$ -dropouts.
" Combining these estimates of halo mass with the observed luminosities allow for an estimate of a mass-to-light ratio, which can constrain the connection between the galaxies’ observed LF and their clustering (Yangetal.2003;Conroy2006)."," Combining these estimates of halo mass with the observed luminosities allow for an estimate of a mass-to-light ratio, which can constrain the connection between the galaxies' observed LF and their clustering \citep[][]{yang2003a,conroy2006a}."
" Numerous effects can alter the correspondence between bias and halo mass, and we therefore view the presented method as primarily a measure of bias and interpret the halo mass estimates with caution."," Numerous effects can alter the correspondence between bias and halo mass, and we therefore view the presented method as primarily a measure of bias and interpret the halo mass estimates with caution."
 Our simulations assume that the high- galaxy is essentially volume limited and pure., Our simulations assume that the high-redshift galaxy population is essentially volume limited and pure.
" Methods populationto handle incompleteness for LBG samples have already been engineered (e.g.,Bouwensetal.2008;Reddy&Steidel2009).."," Methods to handle incompleteness for LBG samples have already been engineered \citep[e.g.,][]{bouwens2008a,reddy2009a}."
" Contamination at a fractional level of f. will lower the bias by an amount A,~(1-f.) (e.g.,Ouchietal.2004).."," Contamination at a fractional level of $f_{c}$ will lower the bias by an amount $\Deltab\approx(1-f_{c})$ \citep[e.g.,][]{ouchi2004a}."
" Using simulations that include red z~2 galaxy interlopers with f;—0.2 and a bias of b©5.5 (Quadrietal.2007), we find that the measured bias at z~6 is lowered by ~17%."," Using simulations that include red $z\sim2$ galaxy interlopers with $f_{c}=0.2$ and a bias of $b\approx5.5$ \citep{quadri2007a}, we find that the measured bias at $z\sim6$ is lowered by $\sim17\%$."
 This systematic error is comparable to the statistical error on the bias when Λύροιας2200., This systematic error is comparable to the statistical error on the bias when $\Nfields\gtrsim200$.
" Lastly, not every parallel field will be useful for measuring high-redshift galaxy counts owing, e.g., to possible bright star or Galactic reddening contamination, and the efficiency of the presented method may be correspondingly decreased."," Lastly, not every parallel field will be useful for measuring high-redshift galaxy counts owing, e.g., to possible bright star or Galactic reddening contamination, and the efficiency of the presented method may be correspondingly decreased."
" Ithank the anonymous referee for constructive suggestions, as well as Richard Ellis, Chuck Steidel, Masami Ouchi, and Haojing Yan for helpful discussions."," I thank the anonymous referee for constructive suggestions, as well as Richard Ellis, Chuck Steidel, Masami Ouchi, and Haojing Yan for helpful discussions."
" I am supported by a Hubble Fellowship grant, program number HST-HF- provided by NASA from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."," I am supported by a Hubble Fellowship grant, program number HST-HF-51262.01-A provided by NASA from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
Foundation (NSF) for funding under eranut ASTOs-35131.,Foundation (NSF) for funding under grant AST08-35734.
 We acknowledge partial support from a Frouticr erant of Heidelberg Uuiversity funded bv the Corman Excellence Initiative aud from the GermanForschung via the ASTRONET project STAR FORMAT (eraut 05A09VITA)., We acknowledge partial support from a Frontier grant of Heidelberg University funded by the German Excellence Initiative and from the German via the ASTRONET project STAR FORMAT (grant 05A09VHA).
 NL.-NLALL. thanks the Institut fiur Theoretische Astrophysil der Cniversitatt Teidelbere for hospitality., M.-M.M.L. thanks the Institut fürr Theoretische Astrophysik der Universitätt Heidelberg for hospitality.
 R.S.I. also thanks the KIPAC at Stanford Wiiversity aud the Department of Astronomy and Astroplysics at the University of California at Santa Cruz for their wari hospitality during a sabbatical stay iu spring 2010., R.S.K. also thanks the KIPAC at Stanford University and the Department of Astronomy and Astrophysics at the University of California at Santa Cruz for their warm hospitality during a sabbatical stay in spring 2010.
 We acknowledge computing tine at the Leibuiz-BRecheuzeutrun i Carchinge (σπα). tle NSE-supported Texas Advanced Computing Center (USA). and at μοι Supercomputing Centre (Conuauy).," We acknowledge computing time at the Leibniz-Rechenzentrum in Garching (Germany), the NSF-supported Texas Advanced Computing Center (USA), and at Jüllich Supercomputing Centre (Germany)."
 The FLASID code was iu part developed by the DOE-supported Alliances Center for Astrophysical Thermounclear Flashes (ASCT) at the Whiversity of Chicago., The FLASH code was in part developed by the DOE-supported Alliances Center for Astrophysical Thermonuclear Flashes (ASCI) at the University of Chicago.
 We thank the anonvinous referee for the useful cohunents. which helped to improve the paper.," We thank the anonymous referee for the useful comments, which helped to improve the paper."
he density of the core of a cluster evolves during its light crossing time. then the gravitational bIueshift imposed on a CAIBR. photon as it falls into the cluster is not matched ww the gravitational redshift imposed as it climbs out. and so the spectrum of the CAIBR. is modified.,"the density of the core of a cluster evolves during its light crossing time, then the gravitational blueshift imposed on a CMBR photon as it falls into the cluster is not matched by the gravitational redshift imposed as it climbs out, and so the spectrum of the CMBR is modified."
 Quilis. LhaikzmarkmainBoclyEndl234>mainBodyStart1235neez Saez (1995) have calculated the largest possible signal due ο this effect.," Quilis, Ib\'a\\tikzmark{mainBodyEnd1234}\~\tikzmark{mainBodyStart1235}neez Saez (1995) have calculated the largest possible signal due to this effect."
 Gravitational lensing of primorcial anisotropics in the CMDRBR. would. lead το a spatiallv-correlated millimetre/submillimetre-wave backeround signal on scales similar to the Einstein radius of the lensing cluster 11)., Gravitational lensing of primordial anisotropies in the CMBR would lead to a spatially-correlated millimetre/submillimetre-wave background signal on scales similar to the Einstein radius of the lensing cluster 1).
 The frequencev-dependent surface. brightness of the SZ increment to the CMDI intensity in the core of a cluster with y=L4.10.! should peak at about μον ? at a wavelength of about jm 110)., The frequency-dependent surface brightness of the SZ increment to the CMBR intensity in the core of a cluster with $y=1.4\times10^{-4}$ should peak at about $\mu$ $^{-2}$ at a wavelength of about $\mu$ m 1).
 In comparison. the largest values of surface brightness of resolved lensed images and galaxies in the cluster at ος=0.171 that was moceled in 11 were about GO and μον aaresce7 respectively.," In comparison, the largest values of surface brightness of resolved lensed images and galaxies in the cluster at $z_{\rm c}=0.171$ that was modeled in 1 were about 60 and $\mu$ $^{-2}$ respectively."
 The RS cllect ancl primordial CAIBR. anisotropics are expected to produce much smaller. surface. brightnesses of order 300 ancl respectively at their maximum intensities 11)., The RS effect and primordial CMBR anisotropies are expected to produce much smaller surface brightnesses of order 300 and $^{-2}$ respectively at their maximum intensities 1).
 The millimetre/submillimetro-wave Dux densities of any clusty star-forming galaxies within a cluster are expected to be considerably smaller than those of the lensecl images of distant galaxies unless the galaxies are in a cluster at a low redshift. that is if z; is significantly less than 0.1 1).," The millimetre/submillimetre-wave flux densities of any dusty star-forming galaxies within a cluster are expected to be considerably smaller than those of the lensed images of distant galaxies unless the galaxies are in a cluster at a low redshift, that is if $z_{\rm c}$ is significantly less than 0.1 1)."
 The spectral energy. distributions of distant. lensed. images should be Hatter than those of cluster galaxies. ancl so at higher frequencies the cluster galaxies should be relatively brighter as compared with the lensecl images: however. the lensecl images are still expected to be brighter than the cluster galaxies throughout the millimetre/submillimetre waveband.," The spectral energy distributions of distant lensed images should be flatter than those of cluster galaxies, and so at higher frequencies the cluster galaxies should be relatively brighter as compared with the lensed images; however, the lensed images are still expected to be brighter than the cluster galaxies throughout the millimetre/submillimetre waveband."
 The typical ux densities of the lensed images should. approximately follow the envelope to the observed limits to the background radiation intensity shown in Fig., The typical flux densities of the lensed images should approximately follow the envelope to the observed limits to the background radiation intensity shown in Fig.
 2., 2.
 Lensed images are expected to dominate the appearance of clusters on arcsecond scales: however. on larger angular scales the surface. brightness distribution of the SZ cllect becomes more significant. and the SZ signal is expected to dominate in observations on arcminute scales at jim. Source confusion is the uncertainty introduced. into the results of an observation by the Dux densities of unresolved sources that lie at. unknown positions in the observing beam.," Lensed images are expected to dominate the appearance of clusters on arcsecond scales; however, on larger angular scales the surface brightness distribution of the SZ effect becomes more significant, and the SZ signal is expected to dominate in observations on arcminute scales at $\mu$ m. Source confusion is the uncertainty introduced into the results of an observation by the flux densities of unresolved sources that lie at unknown positions in the observing beam."
 Source confusion due to extrapolated: and. inferred populations of distant clusty galaxies has been estimated by Franceschini et al. (, Source confusion due to extrapolated and inferred populations of distant dusty galaxies has been estimated by Franceschini et al. (
"1989. 1991). Ποσο Beichman (1990). ""olfolatti et αἱ. (","1989, 1991), Helou Beichman (1990), Toffolatti et al. ("
1995) and. Gawiser Smoot (1997).,1995) and Gawiser Smoot (1997).
 BIS have recently made the first estimates to be based on the results of direct. submillimetre-wave: observations., BIS have recently made the first estimates to be based on the results of direct submillimetre-wave observations.
 Fischer Lange (1993) discussed the cllects of unlensed source confusion on observations of the SZ cllect. and estimated that the confusion noise at a wavelength. of yma in a l-arcmin observing beam should be about 20 times smaller than the signal of the SZ effect [rom a cluster with y=10.3.," Fischer Lange (1993) discussed the effects of unlensed source confusion on observations of the SZ effect, and estimated that the confusion noise at a wavelength of $\mu$ m in a 1-arcmin observing beam should be about 20 times smaller than the signal of the SZ effect from a cluster with $y=10^{-4}$."
 BIS estimate that source confusion at ji is actually about a factor of 7 times more severe than this earlier value., BIS estimate that source confusion at $\mu$ m is actually about a factor of 7 times more severe than this earlier value.
 Confusion. noise in the field of a cluster should be increased by gravitational lensing. because both the mean separation and lux densitv of bright lensed. images are expected. to be larger than those of unlensed background ealaxies.," Confusion noise in the field of a cluster should be increased by gravitational lensing, because both the mean separation and flux density of bright lensed images are expected to be larger than those of unlensed background galaxies."
 In the following sections we estimate the size of this increase and discuss its effects on observations of the SZ ellect., In the following sections we estimate the size of this increase and discuss its effects on observations of the SZ effect.
 In order to make an accurate measurement of the SZ ellect. the Lux densities of confusing sources may need to be subtracted from the observed signal.," In order to make an accurate measurement of the SZ effect, the flux densities of confusing sources may need to be subtracted from the observed signal."
 This is usually achieved by observing the cluster to the same limiting flux density but at a finer angular resolution. in order to detect. the confusing galaxies while resolving out the SZ signal.," This is usually achieved by observing the cluster to the same limiting flux density but at a finer angular resolution, in order to detect the confusing galaxies while resolving out the SZ signal."
 Loeb Relregier (1997). have shown that this procedure coulc introduce a subtle bias into measurements of the SZ elfect., Loeb Refregier (1997) have shown that this procedure could introduce a subtle bias into measurements of the SZ effect.
 Beeause lensed images are typically brighter than unlensec ealaxies. a larger [fraction of their total number can be detected at any fux density limit as compared with unlensec galaxies.," Because lensed images are typically brighter than unlensed galaxies, a larger fraction of their total number can be detected at any flux density limit as compared with unlensed galaxies."
 The images are expected to be more strongly magnilfied in the inner few arcminutes of the field of a cluster. and so source subtraction is expected to be more ellicien there as compared with the outer regions of the field.," The images are expected to be more strongly magnified in the inner few arcminutes of the field of a cluster, and so source subtraction is expected to be more efficient there as compared with the outer regions of the field."
 It is this dillerence in elliciencies that introduces the systematic error into the SZ measurement., It is this difference in efficiencies that introduces the systematic error into the SZ measurement.
 However. this ellect is not of immediate concern in the millimetre/submillimetre waveband. because accurate source subtraction will be impossible in this wavebancl until large interferometer arrays are available (Brown 1996: Downes 1994. 1996).," However, this effect is not of immediate concern in the millimetre/submillimetre waveband, because accurate source subtraction will be impossible in this waveband until large interferometer arrays are available (Brown 1996; Downes 1994, 1996)."
 Phe current accuracy o£ millimetre-wave observations of the SZ ellect is limited by the uncertain Εαν densities of all the confusing ealaxies anc not just by their incomplete subtraction., The current accuracy of millimetre-wave observations of the SZ effect is limited by the uncertain flux densities of all the confusing galaxies and not just by their incomplete subtraction.
 The ellects of gravitational lensing on source confusion was estimated. by combining several dillerent models of ealaxy evolution (Blain Longair 1996) with a model of lensing by clusters (Paper 1)., The effects of gravitational lensing on source confusion was estimated by combining several different models of galaxy evolution (Blain Longair 1996) with a model of lensing by clusters (Paper 1).
 X large set. of. random distributions of background galaxies and their lensed images were derived. and their surface brightness distributions were convolved with Gaussian beams on various scales in order to produce distributions of the tus densities due to both Iensed and unlensecl sources in the beams.," A large set of random distributions of background galaxies and their lensed images were derived, and their surface brightness distributions were convolved with Gaussian beams on various scales in order to produce distributions of the flux densities due to both lensed and unlensed sources in the beams."
 Phe widths of these distributions were then used to predict the source confusion noise expected with and without gravitational lensing., The widths of these distributions were then used to predict the source confusion noise expected with and without gravitational lensing.
redshift the available samples and to conduct a detailed study of the surroundings of the observed compact groups.,redshift the available samples and to conduct a detailed study of the surroundings of the observed compact groups.
 A dedicated search for more distant (up to z~0.2) CGs was started in earnest five to six years ago. with the compilation of a catalogue describing a pilot sample of distant groups drawn from the second digital Palomar Observatory Sky Survey (hereafter DPOSS IL) of Iovino et al.," A dedicated search for more distant (up to $\sim$ 0.2) CGs was started in earnest five to six years ago, with the compilation of a catalogue describing a pilot sample of distant groups drawn from the second digital Palomar Observatory Sky Survey (hereafter DPOSS II) of Iovino et al.,"
 2003. which was later complemented by a catalogue of GCs from SDSS early release (Lee et al..," 2003, which was later complemented by a catalogue of GCs from SDSS early release (Lee et al.,"
 2004). and yet more complete catalogues (de Carvalho et al..," 2004), and yet more complete catalogues (de Carvalho et al.,"
 2005: MeConnachie et al..," 2005; McConnachie et al.,"
 2009. Tago et al..," 2009, Tago et al.,"
 2010)., 2010).
 Most of the information contained in these catalogues is based on photometric data. while spectroscopic redshift information is available for at most two galaxies in each group.," Most of the information contained in these catalogues is based on photometric data, while spectroscopic redshift information is available for at most two galaxies in each group."
 Despite this. the main results that could be drawn from these studies were that distant CGs were very similar to their nearby counterparts.," Despite this, the main results that could be drawn from these studies were that distant CGs were very similar to their nearby counterparts."
 We note. however. that the group environment was never considered in any of these aforementioned Until now. no detailed spectroscopic follow-up has been performed for any distant CGs sample. with the exception of two pilot studies. one by Pompei et al. 2006. for a small sample of DPOSS II CGs. and another by Gutierrez. 2011 for three CGs at z~ 0.3 drawn from the catalogue of MeConnachie et al. (," We note, however, that the group environment was never considered in any of these aforementioned Until now, no detailed spectroscopic follow-up has been performed for any distant CGs sample, with the exception of two pilot studies, one by Pompei et al, 2006, for a small sample of DPOSS II CGs, and another by Gutierrez, 2011 for three CGs at $\sim$ 0.3 drawn from the catalogue of McConnachie et al. ("
2009).,2009).
 This has so far limited any deeper study of the properties of CGs outside our neighbourhood., This has so far limited any deeper study of the properties of CGs outside our neighbourhood.
 To obviate this lack. we started a spectroscopic follow-up campaign on the DPOSS II compact group sample. from 2004 to 2008. observing each member galaxy in 138 candidate compact groups from the DPOSS II catalogue. which was the only large catalogue of distant CGs available when our observations We present in this paper our main results for the whole sample of 138 CGs candidates. deferring to a future paper the discussion of the spectroscopic properties of the member galaxies. the percentage of active galactic nuclei. and the presence of anemic spirals in compact The paper is organized as follows: in Section 2. we describe our observations and data reduction. in Section 3 we present our results. and in Section 4 we discuss the possible implications for the evolution of CGs.," To obviate this lack, we started a spectroscopic follow-up campaign on the DPOSS II compact group sample, from 2004 to 2008, observing each member galaxy in 138 candidate compact groups from the DPOSS II catalogue, which was the only large catalogue of distant CGs available when our observations We present in this paper our main results for the whole sample of 138 CGs candidates, deferring to a future paper the discussion of the spectroscopic properties of the member galaxies, the percentage of active galactic nuclei, and the presence of anemic spirals in compact The paper is organized as follows: in Section 2, we describe our observations and data reduction, in Section 3 we present our results, and in Section 4 we discuss the possible implications for the evolution of CGs."
 In Section 5 we provide our conclusions., In Section 5 we provide our conclusions.
 The sample was selected from the DPOSS II compact group catalogue (lovino et al. 2003.," The sample was selected from the DPOSS II compact group catalogue (Iovino et al. \cite{iovino03},"
 de Carvalho et al.2005)) depending on the allocated observing windows., de Carvalho et \cite{reina05}) ) depending on the allocated observing windows.
 The most comprehensive coverage was between 09< RA xI7h and -1< DEC x+15°. but a few candidates at other coordinates were also observed.," The most comprehensive coverage was between 09$\le$ RA $\le$ 17h and $^{o} \le$ DEC $\le +15^{o}$, but a few candidates at other coordinates were also observed."
 This sample is representative of the DPOSS CGs catalogue. but is by no means complete in either magnitude or redshift.," This sample is representative of the DPOSS CGs catalogue, but is by no means complete in either magnitude or redshift."
 The observations and data reduction were carried out m the same way as described in Pompei et al..," The observations and data reduction were carried out in the same way as described in Pompei et al.,"
" 2006. hereafter Paper I. and we describe them here briefly for completeness All the data were obtained with the 3.58m New Technology Telescope (NTT) and the ESO Multi Mode Instrument (hereafter EMMI) in spectroscopic mode in the red arm.equipped with grism #2 and a slit of 1.5"". under clear/thin cirrus conditions and grey time."," 2006, hereafter Paper I, and we describe them here briefly for completeness All the data were obtained with the 3.58m New Technology Telescope (NTT) and the ESO Multi Mode Instrument (hereafter EMMI) in spectroscopic mode in the red arm,equipped with grism $\#$ 2 and a slit of $\arcsec$, under clear/thin cirrus conditions and grey time."
" The MIT/LL red arm detector. a mosaic of two CCDs 2048 x 4096. was binned by two in both the spatial and spectral directions. with a resulting dispersion. of 3.56A/pix. a spatial scale of 0.33""/pix. an instrumental resolution of 322kms7!.. and a wavelength coverage from 3800 A to 9200 A."," The MIT/LL red arm detector, a mosaic of two CCDs 2048 x 4096, was binned by two in both the spatial and spectral directions, with a resulting dispersion of $\AA$ /pix, a spatial scale of $\arcsec$ /pix, an instrumental resolution of 322, and a wavelength coverage from 3800 $\AA$ to 9200 $\AA$."
 When possible. two or nore galaxies were placed together in the slit. whose position angle had been constrained by the location of galaxies in the sky and thus almost never coincided with the parallactic angle.," When possible, two or more galaxies were placed together in the slit, whose position angle had been constrained by the location of galaxies in the sky and thus almost never coincided with the parallactic angle."
 Exposure times varied from 720s to 1200s per spectrum. and two spectra were taken for each galaxy to ensure reliable cosmic ray subtraction.," Exposure times varied from 720s to 1200s per spectrum, and two spectra were taken for each galaxy to ensure reliable cosmic ray subtraction."
 When the weather conditions allowed it. spectrophotometric standard stars were observed during the night. to flux calibrate the reduced spectra.," When the weather conditions allowed it, spectrophotometric standard stars were observed during the night, to flux calibrate the reduced spectra."
 Standard data reduction was performed using the MIDAS data reduction and our own seripts., Standard data reduction was performed using the MIDAS data reduction and our own scripts.
 Wavelength calibration was applied to the two-dimensional (2D) spectra and an upper limit of 0.16A was found for the rms of the wavelength The two one-dimensional spectra available for each galaxy were averaged together at the end of the reduction. giving an average signal-to-noise ratio (S/N) of ~ 30 (grey time) or  10-15 (almost full moon) per resolution element at 60004. Flux calibration was possible for 805€ of the nights. with an average error of 10-15%.," Wavelength calibration was applied to the two-dimensional (2D) spectra and an upper limit of $\AA$ was found for the rms of the wavelength The two one-dimensional spectra available for each galaxy were averaged together at the end of the reduction, giving an average signal-to-noise ratio (S/N) of $\sim$ 30 (grey time) or $\sim$ 10-15 (almost full moon) per resolution element at $\AA$ Flux calibration was possible for $\%$ of the nights, with an average error of $\%$."
 For the other nights. the conditions were too variable to obtain a reliable Radial velocity standards from the Andersen et al. (," For the other nights, the conditions were too variable to obtain a reliable Radial velocity standards from the Andersen et al. ("
1985) paper were observed with the same instrumental set-up used for the target galaxies: in addition to this. we also used galaxy templates with known spectral characteristies and heliocentric velocity available from the literature. te. M32. NGC 7507. and NGC 4111.,"1985) paper were observed with the same instrumental set-up used for the target galaxies; in addition to this, we also used galaxy templates with known spectral characteristics and heliocentric velocity available from the literature, i.e. M32, NGC 7507, and NGC 4111."
 The technique is similar to the one used in Paper I. and again we briefly mention the points that are most important to understand the current data The packages and were used to measure the galaxy redshifts by means of a cross-correlation," The technique is similar to the one used in Paper I, and again we briefly mention the points that are most important to understand the current data The packages and were used to measure the galaxy redshifts by means of a cross-correlation"
bv the random distribution are considered.,by the random distribution are considered.
 Laciviclual subsources radiate in the field. line direction with a finite angular spread., Individual subsources radiate in the field line direction with a finite angular spread.
 The observed. single pulse is considered. as superposition of many subsources that naturally provide he microstructure., The observed single pulse is considered as superposition of many subsources that naturally provide the microstructure.
 Their random. distribution. leads to luctuations in intensity. [rom pulse to. pulse., Their random distribution leads to fluctuations in intensity from pulse to pulse.
 To. mocel emission from multiple sources. one assumes that the radio emisison is in the plasma natural modes that can escape he pulsar magnetosphere (c.g. Melrose. 2000).," To model emission from multiple sources, one assumes that the radio emisison is in the plasma natural modes that can escape the pulsar magnetosphere (e.g. Melrose 2000)."
 Phe radio enission propagates away [rom the emission region with he polarization being that of the natural modes up to he polarization limitingὃν regiono (PLR). ονομά which the »olarization is no longer alfected by the plasma and is frozen o its value at PLR (Melrose Stoneham LOTT: Darnard Arons 1986).," The radio emission propagates away from the emission region with the polarization being that of the natural modes up to the polarization limiting region (PLR), beyond which the polarization is no longer affected by the plasma and is frozen to its value at PLR (Melrose Stoneham 1977; Barnard Arons 1986)."
 Specifically. single pulses are simulated numerically with the polarization properties derived. from he local plasma clispersion at the PLR.," Specifically, single pulses are simulated numerically with the polarization properties derived from the local plasma dispersion at the PLR."
 The multiple subsource model is described in details in Sec., The multiple subsource model is described in details in Sec.
 2., 2.
 Numerical simulation of single pulses and the implications for the interpretation of the microstructure and fluctuations in intensity are discussed in Sec., Numerical simulation of single pulses and the implications for the interpretation of the microstructure and fluctuations in intensity are discussed in Sec.
 3., 3.
 Conclusions are elven in Sec., Conclusions are given in Sec.
 4., 4.
 Single pulse emission. is modeled. as superposition of emission from a random cistribution of subsources in the emission region., Single pulse emission is modeled as superposition of emission from a random distribution of subsources in the emission region.
 A nonstationary pair cascade above the PC produces a nonsteads. inhomogeneous pulsar. plasma.," A nonstationary pair cascade above the PC produces a nonsteady, inhomogeneous pulsar plasma."
 The corresponding source is then highly inhomogeneous and nonstationary., The corresponding source is then highly inhomogeneous and nonstationary.
 One may model such source in terms of a distribution of multiple subsources., One may model such source in terms of a distribution of multiple subsources.
 Polarized radio emission can be completely described by the Stokes parameters (4. U. Q. V).," Polarized radio emission can be completely described by the Stokes parameters $I$ , $U$ , $Q$, $V$ )."
 Assuming that these subsources are not. phase related. the Stokes parameters can be written as à sum of those from individual emitters. where £* are the intensities of radiation from the ith source in two orthogonal modes d. x; is the position angle (PA) of à rav originating from the ;/th subsource in the observer's direction. and £ and £ are the degree of linear and circular polarization. given by with ZY the polarization ellipticity of the | mode.," Assuming that these subsources are not phase related, the Stokes parameters can be written as a sum of those from individual emitters, where $I^\pm_i$ are the intensities of radiation from the $i$ th source in two orthogonal modes $\pm$, $\chi_i$ is the position angle (PA) of a ray originating from the $i$ th subsource in the observer's direction, and $\xi^l$ and $\xi^c$ are the degree of linear and circular polarization, given by with $T$ the polarization ellipticity of the $+$ mode."
" In the observer's framo. racdation from subsources is beamed in the direction of motion of the source with an angular spread NO,= L/L. where Py is the bulk Lorentz factor of the source fas shown in Figure 12)."," In the observer's frame, radiation from subsources is beamed in the direction of motion of the source with an angular spread $\Delta\Omega_0=1/\Gamma_s$ where $\Gamma_s$ is the bulk Lorentz factor of the source (as shown in Figure \ref{fig:ray}) )."
 Both aberration and refraction can change the beaming direction. substantially Dlaskiewicz. Cordes Wasserman 1991: Petrova 2000: Fussell Luo 2004).," Both aberration and refraction can change the beaming direction substantially (Blaskiewicz, Cordes Wasserman 1991; Petrova 2000; Fussell Luo 2004)."
 The latter elfect. can also cause the two modes to separate (Melrose Stoneham 1977)., The latter effect can also cause the two modes to separate (Melrose Stoneham 1977).
 ALL these ellects ave ignored in the following discussion., All these effects are ignored in the following discussion.
" Phe angular distribution of the intensity of the individual subsource is written as where k is the direction of the line of sight. 4E is the central intensity of the beam. and 8,5 is the propagation angle with respect to the field line direction b."," The angular distribution of the intensity of the individual subsource is written as where $\hat{\bf k}$ is the direction of the line of sight, $I^\pm_{0i}$ is the central intensity of the beam and $\theta_{kb}$ is the propagation angle with respect to the field line direction $\hat{\bf b}$."
" One may express 6,5 in terms of the polar angles of k and b given respectively by (65.04) and (65.605)."," One may express $\theta_{kb}$ in terms of the polar angles of $\hat{\bf k}$ and $\hat{\bf b}$ given respectively by $(\theta_k,\phi_k)$ and $(\theta_b,\phi_b)$."
" In the relativistic limit Pyὃν1. it is convenient to use the small angle approximation [6—6,|« land Jo.ὧν&1."," In the relativistic limit $\Gamma_s\gg1$, it is convenient to use the small angle approximation $|\theta_k-\theta_b|\ll1$ and $|\phi_k-\phi_b|\ll1$."
" Then. the propagation angle in the observers frame is written as Bcx(0.0,37|sin£,(oj,or]2"," Then, the propagation angle in the observer's frame is written as $\theta_{kb}\approx \left[(\theta_k-\theta_b)^2+
\sin\theta_k\sin\theta_b(\phi_k-\phi_b)^2\right]^{1/2}$."
 The distribution of subsources is closely related to that of the pair cascade above the PC., The distribution of subsources is closely related to that of the pair cascade above the PC.
 Phe dominant. pair creation process is single photon decay in the pulsar magnetic field. which is the most cllictent on the field. lines with the smallest. radius of curvature.," The dominant pair creation process is single photon decay in the pulsar magnetic field, which is the most efficient on the field lines with the smallest radius of curvature."
 “Pherefore. pair. production should be peaked on the field lines near the surface boundary subtended by the last open field lines where the electric field is strong and the raclius of field line curvature is small (c.g.Arons Scharlemann1979: Daugherty Harding 1959).," Therefore, pair production should be peaked on the field lines near the surface boundary subtended by the last open field lines where the electric field is strong and the radius of field line curvature is small (e.g.Arons Scharlemann1979; Daugherty Harding 1982)."
 In contrast. there are few pairs created near the magnetic pole where the field line curvature tends to become infinite large.," In contrast, there are few pairs created near the magnetic pole where the field line curvature tends to become infinite large,"
G).,.
 Then. assuming that the last pair of the cascade carries an energy Ej;~20TeV (so that the photon produced through the interaction. with. the CMB caries a typical energy <Γ Τεν). one finds that a magnetic field larger than ~[0715 G isotropizes the low energy cascade. in agreement with the estimates of ?..," Then, assuming that the last pair of the cascade carries an energy $E_{\rm
 fin}\,\sim\,20\,$ TeV (so that the photon produced through the interaction with the CMB carries a typical energy $\lesssim 1\,$ TeV), one finds that a magnetic field larger than $\sim 10^{-12}\,$ G isotropizes the low energy cascade, in agreement with the estimates of \cite{GA05}."
 This situation is modified when one takes into account the inhomogeneous distribution of extra-galactie magnetic fields. as we now discuss.," This situation is modified when one takes into account the inhomogeneous distribution of extra-galactic magnetic fields, as we now discuss."
 Primary cosmic rays. upon traveling through the voids of large scale structure may inject secondary pars which undergo inverse Compton cascades in these unmagnetized regions.," Primary cosmic rays, upon traveling through the voids of large scale structure may inject secondary pairs which undergo inverse Compton cascades in these unmagnetized regions."
 If the field in such regions is smaller than the above 107' G. then the cascade will transmit. its energy in forward <TeV photons.," If the field in such regions is smaller than the above $10^{-12}\,$ G, then the cascade will transmit its energy in forward $\lesssim $ TeV photons."
 Of course. depending on the exact value of B where the cascade ends. the resulting image will be spread by some finite angle.," Of course, depending on the exact value of $B$ where the cascade ends, the resulting image will be spread by some finite angle."
 Since we are interested in sharply peaked images. let us consider a typical angular size 6 and ignore those regions in which the magnetic field is large enough to give a contribution to the Image on a size larger than &.," Since we are interested in sharply peaked images, let us consider a typical angular size $\theta$ and ignore those regions in which the magnetic field is large enough to give a contribution to the image on a size larger than $\theta$."
 For €.«1. the problem remains one-dimensional as before. and one cancompute the total energy injected in inverse Compton cascades within 8. as follows.," For $\theta\ll1$, the problem remains one-dimensional as before, and one cancompute the total energy injected in inverse Compton cascades within $\theta$, as follows."
 The luminosity injected in secondary pairs and photons up to distance d is written χοζωE)., The luminosity injected in secondary pairs and photons up to distance $d$ is written $\chi_e L_{\rm cr}(>E)$.
 Since we are interested in the signatures of ultrahigh cosmic ray sources. we require that E>10 eV; for protons. the energy loss length due pair production moreover increases dramatically as E becomes smaller than 10'? eV. so that the contribution of lower energy particles can be neglected in a first approximation.," Since we are interested in the signatures of ultrahigh cosmic ray sources, we require that $E\geq10^{19}\,$ eV; for protons, the energy loss length due pair production moreover increases dramatically as $E$ becomes smaller than $10^{19}\,$ eV, so that the contribution of lower energy particles can be neglected in a first approximation."
 For photo-pair production. the fraction transfered 18. yi1Gpe of Ley=ἰωδ10eV) up to d~I Gpe.," For photo-pair production, the fraction transfered is $\chi_{e,ee}\simeq d/1\,{\rm
 Gpc}$ of $L_{E,19}=L_{\rm cr}(>10^{19}\,{\rm eV})$ up to $d\sim1\,$ Gpc."
" For pion production, the fraction of energy transfered is roughly ye.3 of ἐς610!eV) in the continuous energy loss approximation."," For pion production, the fraction of energy transfered is roughly $\chi_{e,\pi}\simeq d/100\,{\rm Mpc}$ of $L_{\rm
 cr}(>6\,10^{19}\,{\rm eV})$ in the continuous energy loss approximation."
" At distances I00MpcxdlÓGpe. the fraction y, of Lyio injected into secondary pairs and photons thus ranges from ~0.5 ford=100 Mpe to ~ Latd=1 Gpe: in short. it is expected to be of order unity or slightly less."," At distances $100\,{\rm Mpc}\,\leq\,d\,\leq\,1\,{\rm
 Gpc}$, the fraction $\chi_e$ of $L_{E,19}$ injected into secondary pairs and photons thus ranges from $\sim
 0.5$ for $d=100\,$ Mpc to $\sim1$ at $d=1\,$ Gpc; in short, it is expected to be of order unity or slightly less."
 All the energy injected in this way in sufficiently unmagnetized regions (see below) will be deposited through the inverse Compton cascade in the sub-TeV range. with a typical energy flux dependence «E!. up to some maximal energy Eynay1-10 TeV beyond which the Universe is opaque to gamma rays on the distance scale d (?)..," All the energy injected in this way in sufficiently unmagnetized regions (see below) will be deposited through the inverse Compton cascade in the sub-TeV range, with a typical energy flux dependence $\propto
 E_\gamma^{1/2}$ up to some maximal energy $E_{\gamma,\rm max}\sim
 1-10\,$ TeV beyond which the Universe is opaque to gamma rays on the distance scale $d$ \citep{Ferrigno04}."
" Neglecting any redshift dependence for simplicity. the gamma-ray energy flux per unit energy interval may then be approximated as: where fiu(<B,) denotes the one-dimensional filling factor. i.e. the fraction of the line of sight in which the magnetic field is smaller than the value δρ such that the deflection of the low energy cascade is 4."," Neglecting any redshift dependence for simplicity, the gamma-ray energy flux per unit energy interval may then be approximated as: where $f_{\rm 1d}(<B_\theta)$ denotes the one-dimensional filling factor, i.e. the fraction of the line of sight in which the magnetic field is smaller than the value $B_\theta$ such that the deflection of the low energy cascade is $\theta$."
 For reference. Bj~2x107 G ford~1°," For reference, $B_\theta
\,\sim\,2\times 10^{-14}\,$ G for $\theta\sim1^\circ$."
 In general. one finds in the literature the three-dimensional filling factor fij. but fiq(«Ba)~αςBy) up to a numerical prefactor of order unity that depends on the geometry of the structures.," In general, one finds in the literature the three-dimensional filling factor $f_{\rm
 3d}$, but $f_{\rm 1d}(<B_\theta) \sim f_{\rm 3d}(<B_\theta)$ up to a numerical prefactor of order unity that depends on the geometry of the structures."
 Interestingly enough. the amount of magnetizatior of the voids of large scale structure 1s directly related to the origin of large scale magnetic fields.," Interestingly enough, the amount of magnetization of the voids of large scale structure is directly related to the origin of large scale magnetic fields."
 Obviously. if galactic anc cluster magnetic fields originate from a seed field produced 1 a homogeneous way with a present day strength Bc»107! G. then the above gamma ray flux will be diluted to large angular scales. hence below detection threshold.," Obviously, if galactic and cluster magnetic fields originate from a seed field produced in a homogeneous way with a present day strength $B\gg 10^{-14}\,$ G, then the above gamma ray flux will be diluted to large angular scales, hence below detection threshold."
" However. if the seed field. extrapolated to present day values is much lower than this value. or if most of the magnetic enrichment of the intergalactic medium results from the pollution by star forming galaxies and radio-galaxies. then one should expect fiu(<B,) to be non negligible."," However, if the seed field, extrapolated to present day values is much lower than this value, or if most of the magnetic enrichment of the intergalactic medium results from the pollution by star forming galaxies and radio-galaxies, then one should expect $f_{\rm 1d}(<B_\theta)$ to be non negligible."
 For instance. ? obtain Auc1077G)-O(1) in such models.," For instance, \cite{2009MNRAS.392.1008D} obtain $f_{\rm 3d}(<10^{-14}\,{\rm G})\sim
{\cal O}(1)$ in such models."
" Given the sensitivity of current gamma ray experiments such as H.E.S.S. in the TeV energy range. the inverse Compton cascades might then produce degree-size detectable halos for source luminosities >10""|faq1077Gycd/100Mpe)7 erg/s. For future instruments such as CTA. the detectability condition will be of order L,>10!fice1077Gy]7(4/100Mpc)7 erg/s. The expected flux level remains quite uncertain as it depends mostly on the configuration of the extragalactic magnetic field in the voids. contrarily to the synchrotron signal which is mainly controlled by the source luminosity."," Given the sensitivity of current gamma ray experiments such as H.E.S.S. in the TeV energy range, the inverse Compton cascades might then produce degree-size detectable halos for source luminosities $\gtrsim 10^{43}[f_{\rm 3d}(<10^{-14}\,{\rm G})]^{-1}(d/100\,{\rm
 Mpc})^{-2}\,$ erg/s. For future instruments such as CTA, the detectability condition will be of order $L_{\rm
 s}\gtrsim10^{41}[f_{\rm 3d}(<10^{-14}\,{\rm G})]^{-1}(d/100\,{\rm Mpc})^{-2}\,$ erg/s. The expected flux level remains quite uncertain as it depends mostly on the configuration of the extragalactic magnetic field in the voids, contrarily to the synchrotron signal which is mainly controlled by the source luminosity."
" We note that intergalactic magnetic fields of strength 107 G might be probed through the delay time of the high energy afterglow of gamma-ray bursts (2?) or the GeV emission around blazars (22222)., as discussed earlier."," We note that intergalactic magnetic fields of strength $B\lesssim10^{-15}\,$ G might be probed through the delay time of the high energy afterglow of gamma-ray bursts \citep{1995Natur.374..430P,2008ApJ...682..127I} or the GeV emission around blazars \citep{ACV94,Dai02,dAvezac:2007sg,2009PhRvD..80b3010E,2009PhRvD..80l3012N}, as discussed earlier."
 As the present paper was being refereed. some first estimates on lower bounds for the average magnetic field appeared. giving 10776—10° G. see 2.. 2? and ?..," As the present paper was being refereed, some first estimates on lower bounds for the average magnetic field appeared, giving $B\gtrsim 10^{-16}-10^{-15}$ G, see \cite{Neronov10}, \cite{Ando10} and \cite{Dolag10}."
 In view of the results of section ??.. one may want to consider also the case of mildly powerful but nearby sources.," In view of the results of section \ref{subsection:average}, one may want to consider also the case of mildly powerful but nearby sources."
 One such potential source is Cen A. which has attracted a considerable amount of attention. as it is the nearest radio-galaxy (3.8 Mpe) and more recently. because a fraction of the Pierre Auger events above 60 EeV have clustered within 10—20° of this source.," One such potential source is Cen A, which has attracted a considerable amount of attention, as it is the nearest radio-galaxy (3.8 Mpc) and more recently, because a fraction of the Pierre Auger events above $60\,$ EeV have clustered within $10-20^\circ$ of this source."
 As discussed in detail in 2.. this source is likely too weak to accelerate protons to =>10' eV in steady state. but heavy nuclei might possibly be accelerated up to GZK energies.," As discussed in detail in \cite{LW09}, this source is likely too weak to accelerate protons to $\gtrsim 10^{19}\,$ eV in steady state, but heavy nuclei might possibly be accelerated up to GZK energies."
 The acceleration of protons to ultrahigh energies in flaring episodes of high luminosity has been proposed in ?.., The acceleration of protons to ultrahigh energies in flaring episodes of high luminosity has been proposed in \cite{DA09}.
 It is also important to recall that the apparent clustering mn this direction can be naturally explained thanks to the large concentration of matter in the local Universe. in the direction of Cen A but located further away.," It is also important to recall that the apparent clustering in this direction can be naturally explained thanks to the large concentration of matter in the local Universe, in the direction of Cen A but located further away."
 As discussed in ?.. some of the events detected toward Cen A could also be attributed to rare and powerful bursting sources located in the Cen A host galaxy. such as gamma-ray bursts.," As discussed in \cite{LW09}, , some of the events detected toward Cen A could also be attributed to rare and powerful bursting sources located in the Cen A host galaxy, such as gamma-ray bursts."
 Due to the scattering of the particles on the magnetized lobes. one would not be able to distinguish such bursting sources from a source located in the core of Cen A.," Due to the scattering of the particles on the magnetized lobes, one would not be able to distinguish such bursting sources from a source located in the core of Cen A."
jwe. both a very large volune that coutaiuns mauy cluster-zed halos. and high mass aud spatial resolution o resolve scales of ~LOAtkpe.,"have, both a very large volume that contains many cluster-sized halos, and high mass and spatial resolution to resolve scales of $\sim10\hkpc$."
 In fact. on these very πα scales. it is likely that dark matter subhalos have already been disrupted by tidal forces. while the LRGs. cine snaller aud deuser. have survived.," In fact, on these very small scales, it is likely that dark matter subhalos have already been disrupted by tidal forces, while the LRGs, being smaller and denser, have survived."
 So pure dark uatter simulations may be insufficient to predict these LRG results., So pure dark matter simulations may be insufficient to predict these LRG results.
 Towever. seiui-iualvtie models that include ealaxies and can have arbitrarily Ligh resolution should f able to make these predictions (see ποπ]&White2008 for seimi-analvtie modeling of the verv-iuall scale clustering of lower Iuninosity galaxies).," However, semi-analytic models that include galaxies and can have arbitrarily high resolution should be able to make these predictions (see \citealt{kitzbichler08} for semi-analytic modeling of the very-small scale clustering of lower luminosity galaxies)."
 Iu anv case. our results have nuplicatious for the nodchug of LRC clustering because a standard NEW profile caunot describe the spatial distribution of LRGs on scales Z0.0381Ape.," In any case, our results have implications for the modeling of LRG clustering because a standard NFW profile cannot describe the spatial distribution of LRGs on scales $\lesssim 0.03\hmpc$."
 It would be iuteresting to scc if lower huuinositv galaxies exhibit the same behavior or if this is simply a feature for LRCs., It would be interesting to see if lower luminosity galaxies exhibit the same behavior or if this is simply a feature for LRGs.
 It would also be interesting to see if LRGs maintain their steep deusitv profile at high redshift., It would also be interesting to see if LRGs maintain their steep density profile at high redshift.
 We thank Joanna Duuklev aud David Ποσο for useful discussions and couunents., We thank Joanna Dunkley and David Hogg for useful discussions and comments.
 We thank Zheng Zhene for providing us with his best-fit 2-halo term for LRGs., We thank Zheng Zheng for providing us with his best-fit 2-halo term for LRGs.
usingLST.,using.
 Section 2 describes the data obtained and (he methods used {ο process them., Section 2 describes the data obtained and the methods used to process them.
 Section 3 provides our results., Section 3 provides our results.
 Section 4 discusses the implications of the results. including a prediction of the orbital period of the LAINB. and 85 gives our conclusions.," Section 4 discusses the implications of the results, including a prediction of the orbital period of the LMXB, and 5 gives our conclusions."
 We obtained six observations with the ACIS-I between December 2003 and December 2004 relevant to (his study., We obtained six observations with the ACIS-I between December 2003 and December 2004 relevant to this study.
 The observation identification numbers. dates. pointings. roll angles. ancl exposure times of these observations are given in Table L1..," The observation identification numbers, dates, pointings, roll angles, and exposure times of these observations are given in Table \ref{xobs}."
" Although the observations were taken [or 5 ks each. the effective exposure was reduced by ~20'% because the data were taken in. ""alternating readout mode"" in order to avoid pileup for any bright iransient source."," Although the observations were taken for 5 ks each, the effective exposure was reduced by $\sim$ because the data were taken in “alternating readout mode” in order to avoid pileup for any bright transient source."
 These observations were all reduced. using the software package CIAO v3.1 with the CALDD v2.28., These observations were all reduced using the software package CIAO v3.1 with the CALDB v2.28.
" We created exposure maps for the images using the task and we found and measured positions and fluxes of (he sources in (he images using the CIAO task Each data set detected sources down to (0.310 keV) fluxes of ~6x "" photons 7 1 Lor 0.310 keV luminosities of &-107 erg |! [or a (vpical X-ray binary svstem in M31.", We created exposure maps for the images using the task and we found and measured positions and fluxes of the sources in the images using the CIAO task Each data set detected sources down to (0.3–10 keV) fluxes of $\sim$ $\times$ $^{-6}$ photons $^{-2}$ $^{-1}$ or 0.3–10 keV luminosities of $\sim$ $^{36}$ erg $^{-1}$ for a typical X-ray binary system in M31.
 We aligned the coordinate svstem of the N-ray images with the optical images of the Local Group Survey (LGS: Massey.etal. 2001))., We aligned the coordinate system of the X-ray images with the optical images of the Local Group Survey (LGS; \citealp{massey2001}) ).
 The positions of X-ray sources with known elobular cluster counterparts were aligned wilh the globular cluster centers in the LGS images using the taskcemap., The positions of X-ray sources with known globular cluster counterparts were aligned with the globular cluster centers in the LGS images using the task.
" The errors of this alignment were tvpicallv 70.1"".", The errors of this alignment were typically $\sim$ $''$.
 The precise alignment. errors for the observations used (o determine the A-ray error circle are given in Table 3.., The precise alignment errors for the observations used to determine the X-ray error circle are given in Table \ref{xpos}.
 We checked the positions of all of the sources we detected in each observation wil those, We checked the positions of all of the sources we detected in each observation with those
also be noted that at high metallicities the value of this index changes rapidly with metallicity.,also be noted that at high metallicities the value of this index changes rapidly with metallicity.
 Lt is possible that the seeming Fet66s excess is in [act the result of the lack of calibrated SSP data at metallicities greater than Ve/LI] = 0.5., It is possible that the seeming Fe4668 excess is in fact the result of the lack of calibrated SSP data at metallicities greater than [Fe/H] = 0.5.
 Caution is acvisecl in interpreting this index at hisgh metallicitics., Caution is advised in interpreting this index at high metallicities.
 Cradients were estimated: by least squares fitting for the «pre. Fot66S. Mg» and 11.2 indices plotted versus log radius from the galaxy contre (fig. 1)).," Gradients were estimated by least squares fitting for the $<$ $>$ , Fe4668, $_{2}$ and ${\beta}$ indices plotted versus log radius from the galaxy centre (fig. \ref{grads}) )."
 Data from both sides of the galaxies were included. in these fits., Data from both sides of the galaxies were included in these fits.
 The gradients obtained are given in table 3.., The gradients obtained are given in table \ref{cindices}.
 Central galaxy. regions are mareinally allectec by seeing., Central galaxy regions are marginally affected by seeing.
 However. central index values were included in the gradient estimates due to the limited number of data. points.," However, central index values were included in the gradient estimates due to the limited number of data points."
 As gradients are unallected by a constant oll-set to convert to the Lick svstem. a cirect comparison with elliptical galaxies can be mace.," As gradients are unaffected by a constant off-set to convert to the Lick system, a direct comparison with elliptical galaxies can be made."
 Ciradients. central Mg» values and central velocity clispersions (ay: table 3)) of the bulges are consistent with the correlations reported by Carollo ct al.," Gradients, central $_{2}$ values and central velocity dispersions $\sigma_{0}$; table \ref{cindices}) ) of the bulges are consistent with the correlations reported by Carollo et al."
 (1t093) in elliptical galaxies., \shortcite{CDB93} in elliptical galaxies.
 The bulges show Meg» gradients of similar magnitude to those in elliptical galaxies possessing the same central velocity dispersions and. central Meo values., The bulges show $_{2}$ gradients of similar magnitude to those in elliptical galaxies possessing the same central velocity dispersions and central $_{2}$ values.
 The gradients are also similar in magnitude to those reported. in Sansom. Peace Dodd (1994) for 2 bulges (NGC 3190 and NGC 1023," The gradients are also similar in magnitude to those reported in Sansom, Peace Dodd \shortcite{SPD94} for 2 bulges (NGC 3190 and NGC 1023)."
 The central velocity clispersions and the central index values of our Palomar sample. indicating solar abundance ratios. are also Consistent with the pattern suggested by Worthey (1998) that spheroids with velocity clispersions less than 225 | possess solar “Me/Fe].," The central velocity dispersions and the central index values of our Palomar sample, indicating solar abundance ratios, are also consistent with the pattern suggested by Worthey \shortcite{W98}
 that spheroids with velocity dispersions less than 225 $^{-1}$ possess solar [Mg/Fe]."
 In order to investigate possible ος. we use our galactic chemical evolution (GCE) code cetailed in Sansom Proctor (1998)..," In order to investigate possible SFHs, we use our galactic chemical evolution (GCE) code detailed in Sansom Proctor \shortcite{SP98a}."
 Briclly. the model calculates star formation rate (SER) and the metallicity of the gas at cach time-step throughout the history of he region being moclellec.," Briefly, the model calculates star formation rate (SFR) and the metallicity of the gas at each time-step throughout the history of the region being modelled."
 As all stars. formed. in each time step are of the same age and metallicity. they constitute an SSP.," As all stars formed in each time step are of the same age and metallicity, they constitute an SSP."
 Composite indices are then caleulated as the luminosity weighted sum of indices [rom interpolations between tabulated SSP values of Worthey (1994) which cover populations of age 1.5r to 1v Gwe., Composite indices are then calculated as the luminosity weighted sum of indices from interpolations between tabulated SSP values of Worthey \shortcite{W94} which cover populations of age 1.5 to 17 Gyr.
 The code currently models a single zone allowing for eas in-lall., The code currently models a single zone allowing for gas in-fall.
 Vhe rate of eas inflow can be varied and can be either primordial or enriched to the same level as the gas already in the region (simulating inflow from a neighbouring region with similar SELL)., The rate of gas inflow can be varied and can be either primordial or enriched to the same level as the gas already in the region (simulating inflow from a neighbouring region with similar SFH).
" SETU is assumed to be proportional to some power (a) of the gas mass density (c (p): SER = Cp"". where € is the star formation elficiency +) if asl."," SFR is assumed to be proportional to some power $\alpha$ ) of the gas mass density $\rho$ ); SFR = $\rho^{\alpha}$, where C is the star formation efficiency $^{-1}$ ) if $\alpha$ =1."
 Ixennicutt (1989) shows that the value of à in galactic discs lies between 1: and , Kennicutt \shortcite{K89} shows that the value of $\alpha$ in galactic discs lies between 1 and 2.
In all models. presented. here we assume a=) and inflow is enriched., In all models presented here we assume $\alpha$ =1 and inflow is enriched.
 Ehe dillerences in indices caused by assuming à = 2 were shown to be small and make no significantdifference to the conclusions for the models presented here., The differences in indices caused by assuming $\alpha$ = 2 were shown to be small and make no significantdifference to the conclusions for the models presented here.
 The code permits the modification of the star formation cllicteney ancl inflow rates at (vo points, The code permits the modification of the star formation efficiency and inflow rates at two points
within the DEEP2 survey area.,within the DEEP2 survey area.
" By comparing to the average dark matter halo bias of the galaxy sample, obtained through a calculation of the galaxy auto-correlation, we estimate the typical bias of the absorbers and thereby quantify their average halo mass."," By comparing to the average dark matter halo bias of the galaxy sample, obtained through a calculation of the galaxy auto-correlation, we estimate the typical bias of the absorbers and thereby quantify their average halo mass."
 We also measure the covering fraction of W2?7°S>0.6A absorption for the DEEP2 galaxy sample and compare with previous results., We also measure the covering fraction of $^{\lambda2796}_{r}\ga$ absorption for the DEEP2 galaxy sample and compare with previous results.
 A description of the data is given in Section 2., A description of the data is given in Section 2.
" The methodology we apply for the analysis is described in Section 3, and results are presented in Sections 4-7."," The methodology we apply for the analysis is described in Section 3, and results are presented in Sections 4-7."
" A discussion of the results is given in Section 8, and a summary of important findings is presented in Section 9."," A discussion of the results is given in Section 8, and a summary of important findings is presented in Section 9."
" Throughout this paper, we assume a flat A-dominated CDM cosmology with Ωμ,=0.27, Ho=73 km s! Mpc'!, and os—0.8 unless otherwise stated."," Throughout this paper, we assume a flat $\Lambda$ –dominated CDM cosmology with $\Omega_m=0.27$, $H_0=73$ km $^{-1} $ $^{-1}$, and $\sigma_8=0.8$ unless otherwise stated."
 'The quasar spectra used in this analysis are drawn from the Sloan Digital Sky Survey Seventh Data Release (SDSSDR7Abazajianetal. 2009)., The quasar spectra used in this analysis are drawn from the Sloan Digital Sky Survey Seventh Data Release \citep[SDSS DR7][]{A09}.
. The SDSS DR7 Quasar catalogue (Schneideretal.2010) contains 88 quasars in regions of overlap with the DEEP2 survey., The SDSS DR7 Quasar catalogue \citep{Schneider10} contains 88 quasars in regions of overlap with the DEEP2 survey.
" Each of the quasars in the DEEP2 survey area was run through an automated pipeline that detects strongMgII absorption systems with high precision (for full description, see York et al."," Each of the quasars in the DEEP2 survey area was run through an automated pipeline that detects strong absorption systems with high precision (for full description, see York et al."
 2011)., 2011).
" In brief, this pipeline first searches for absorption features in quasar spectra."," In brief, this pipeline first searches for absorption features in quasar spectra."
" For each detected absorption feature, a Gaussian profile is fit to the normalised flux to extract precise centroid and equivalent width measurements."," For each detected absorption feature, a Gaussian profile is fit to the normalised flux to extract precise centroid and equivalent width measurements."
" In order to identify the ion and redshift corresponding to each measured absorption line, a system-finding algorithm then identifies pairs of 4σ absorber detections matching the wavelength separation expected for at a given redshift."," In order to identify the ion and redshift corresponding to each measured absorption line, a system-finding algorithm then identifies pairs of $\sigma$ absorber detections matching the wavelength separation expected for at a given redshift."
" The reliability of the identification is further quantified for each detected absorption system, taking into account the doublet ratio measured forII, the number of additional ions matched in absorption at the same redshift, and any blending with other absorption features identified at another redshift."," The reliability of the identification is further quantified for each detected absorption system, taking into account the doublet ratio measured for, the number of additional ions matched in absorption at the same redshift, and any blending with other absorption features identified at another redshift."
" Due to the magnitude-limited design of the SDSS survey, the spectral signal-to-noise ratio (SNR) correlates with the optical apparent magnitude of each quasar."," Due to the magnitude-limited design of the SDSS survey, the spectral signal-to-noise ratio (SNR) correlates with the optical apparent magnitude of each quasar."
" The completeness of the SDSS absorption line data for any object thereby varies simultaneously with absorber equivalent width and quasar magnitude, such that the fainter the quasar, the poorer the spectral SNR, and thus the larger the minimum absorber equivalent width for detection at the required 4σ limit."," The completeness of the SDSS absorption line data for any object thereby varies simultaneously with absorber equivalent width and quasar magnitude, such that the fainter the quasar, the poorer the spectral SNR, and thus the larger the minimum absorber equivalent width for detection at the required $\sigma$ limit."
" A full description of the detection completeness is provided in Yorketal.(2011),, but generally z~1 absorption systems with W2?7%>0.6A are detected with 29096 completeness in SDSS quasars with m; «20."," A full description of the detection completeness is provided in \citet{York11}, but generally $\sim$ 1 absorption systems with $_{r}^{\lambda2796}$$>$ are detected with $>$ completeness in SDSS quasars with $_{i}\leq$ 20."
" Because we include in our sample quasars with m; 220, all candidate absorption systems were verified by visual inspection to rule out any false detections due to artifacts such as problematic continuum subtraction around narrow emission lines or noise spikes."," Because we include in our sample quasars with $_{i}>$ 20, all candidate absorption systems were verified by visual inspection to rule out any false detections due to artifacts such as problematic continuum subtraction around narrow emission lines or noise spikes."
" In all, we find 21 absorption systems with W2779»0.6À and a velocity separation of v 515.000 km s! in the quasar rest frame."," In all, we find 21 absorption systems with $_{r}^{\lambda2796}$$\geq$ and a velocity separation of $v>$ 15,000 km $^{-1}$ in the quasar rest frame."
 This velocity difference is generally sufficient to ensure that the absorbers are unambiguously unassociated with the background quasar and thus originating in foreground galactic environments (Wildetal.2008)., This velocity difference is generally sufficient to ensure that the absorbers are unambiguously unassociated with the background quasar and thus originating in foreground galactic environments \citep{Wild08}.
". The mean rest-frame equivalent width of the 18 absorbers extracted from spectra with m;<20 is -W?965—119À, compared to <W2?7%>=1.88A for the remaining 3 absorbers drawn from fainter quasars."," The mean rest-frame equivalent width of the 18 absorbers extracted from spectra with $_{i}\leq$ 20 is $<$$W_{r}^{\lambda2796}$$>$, compared to $<$$W_{r}^{\lambda2796}$$>$ for the remaining 3 absorbers drawn from fainter quasars."
" While— we expect that the three included spectra with m;220 are only complete to W2?796—1.0À., these data represent a small fraction of the selected sample."," While we expect that the three included spectra with $_{i}>$ 20 are only complete to $_{r}^{\lambda2796}$, these data represent a small fraction of the selected sample."
 Additional details of the absorber detections are provided in Table 1., Additional details of the absorber detections are provided in Table 1.
" The galaxies used in this analysis are taken from the DEEP2 Galaxy Redshift Survey (Davisetal.2003),, which obtained spectra for ~32,000 galaxies using the DEIMOS spectrograph (Faberetal.2003) on the Keck II telescope."," The galaxies used in this analysis are taken from the DEEP2 Galaxy Redshift Survey \citep{Davis03}, which obtained spectra for $\sim$ 32,000 galaxies using the DEIMOS spectrograph \citep{Faber03} on the Keck II telescope."
" The DEEP2 survey observed galaxies in the redshift range 0.7<z<1.45, with a limiting magnitude of Rap=24.1."," The DEEP2 survey observed galaxies in the redshift range $<$ $<$ 1.45, with a limiting magnitude of $R_{AB}$ =24.1."
" A magnitude cut of 18.5<Rag24.1, in combination with a BRI color cut, was applied to select galaxies in the redshift range of interest from CFHT CFH12K imaging (Coiletal.2004b)."," A magnitude cut of $18.5<R_{AB}<24.1$, in combination with a BRI color cut, was applied to select galaxies in the redshift range of interest from CFHT CFH12K imaging \citep{Coil04b}."
". The survey geometry consists of four spatially separated fields, which together cover 3 deg? of the sky."," The survey geometry consists of four spatially separated fields, which together cover 3 $^{2}$ of the sky."
" Additional details of the survey observations and data reduction are available in Davisetal.(2003),, and Coilet (2004a)."," Additional details of the survey observations and data reduction are available in \citet{Davis03}, and \citet{Coil04a}."
". The DEEP2 galaxy catalog is a flux-limited sample with <mg >=23.3 and a mean rest-frame B-band luminosity «Mg >=-20.1, which has been used extensively for clusteringmeasurements at z~1."," The DEEP2 galaxy catalog is a flux-limited sample with $<m_{R}>$ =23.3 and a mean rest-frame B-band luminosity $<M_{B}>$ =-20.1, which has been used extensively for clusteringmeasurements at $\sim$ 1."
" For this work, we apply the exact sample used by Coiletal.(2007) to measure the quasar-galaxy cross-correlation."," For this work, we apply the exact sample used by \citet{Coil07} to measure the quasar-galaxy cross-correlation."
towards shallow density profiles aud rounder shapes.,towards shallow density profiles and rounder shapes.
 We have also tested the impact of eravitational softening on our simulations., We have also tested the impact of gravitational softening on our simulations.
 We performed two siuulations. identical to the standardresolution. (06 halo but with softening increased/decreased by a factor of two.," We performed two simulations, identical to the standard–resolution, $0\sigma$ halo but with softening increased/decreased by a factor of two."
 Decreasing the softening leusth docs uot noticeably change any muportaut quantities. (density. velocity dispersion. axis ratios. phasespace density).," Decreasing the softening length does not noticeably change any important quantities (density, velocity dispersion, axis ratios, phase–space density)."
 Tucreasing the softening length results in a shallower core density profile. due to the decreased gravitational force at σα]. iuter-particle distances.," Increasing the softening length results in a shallower core density profile, due to the decreased gravitational force at small inter-particle distances."
 The phasespace deusity xofile is similarly affected by the softening changes: jowever. the auisotropy profile shows uo difference when 1e softening leugths are changed.," The phase–space density profile is similarly affected by the softening changes; however, the anisotropy profile shows no difference when the softening lengths are changed."
 These effects are ouly roticeable within a radius of 0.057599. the same radius within which we may be affected by resolution Inaitatious 1.05 brogy).," These effects are only noticeable within a radius of $0.05 r_{200}$, the same radius within which we may be affected by resolution limitations $0.054
r_{200}$ )."
 We conclude that our results are robust ονομα a radius of 0.057599., We conclude that our results are robust beyond a radius of $0.05 r_{200}$.
 Plus distance corresponds o 10 kpe for a Milkv Waysize halo., This distance corresponds to 10 kpc for a Milky Way–size halo.
 As the halo collapses. we measure the density profile. velocity dispersion tensor. and the axis ratios as a function of radius at cach time step.," As the halo collapses, we measure the density profile, velocity dispersion tensor, and the axis ratios as a function of radius at each time step."
 We expect to see higher central deusities aud larger radial velocity dispersions with increasing tine., We expect to see higher central densities and larger radial velocity dispersions with increasing time.
 However. if the ROI is present we also expect to see deviations from spherical sviunetry develop iu regions with laree radial velocity dispersions.," However, if the ROI is present we also expect to see deviations from spherical symmetry develop in regions with large radial velocity dispersions."
 These deviations would be accompanied by Increasing tangential velocities due to the additional torques., These deviations would be accompanied by increasing tangential velocities due to the additional torques.
 The most obvious evidence for the ROT is the erowth of triaxialitv iu the shape of the halo., The most obvious evidence for the ROI is the growth of triaxiality in the shape of the halo.
 We see this vchavior clearly iu our simulations., We see this behavior clearly in our simulations.
 We measure axis ratios as a function of radius by dividing the halo into 196 equal particle-unuuber radial bius based ou xwticle density. and then calculating the monmenut-of- tensor for cach biu (see Babul&Starkiman(1992) or ai application of this method).," We measure axis ratios as a function of radius by dividing the halo into 196 equal particle-number radial bins based on particle density, and then calculating the moment-of-inertia tensor for each bin (see \citet{babul92} for an application of this method)."
 The square root of the ratios of the principal moments of iertia in the directions of the three axes of svuuuetry give the axis ratios of he particles withiu each radial bin., The square root of the ratios of the principal moments of inertia in the directions of the three axes of symmetry give the axis ratios of the particles within each radial bin.
 This method more accurately tracks the racial variations of the axis ratios han methods which analyze all particles interior to a eiven radius., This method more accurately tracks the radial variations of the axis ratios than methods which analyze all particles interior to a given radius.
 This latter methocl is always dominated by he shape of the inner halo. where the majority of the virialized halo's mass is located. aud thus is inscusitive to variations in the axis ratios at large radii.," This latter method is always dominated by the shape of the inner halo, where the majority of the virialized halo's mass is located, and thus is insensitive to variations in the axis ratios at large radii."
 After 0.6 Co. the first reaches the center leading to a large radial velocity dispersion.," After 0.6 Gyr, the first reaches the center leading to a large radial velocity dispersion."
 Tn response. a triaxial structure develops iu the core of the halo.," In response, a triaxial structure develops in the core of the halo."
 The resulting axis ratio profiles can be seen in Fieure 2.. which portravs the evolution of the axis ratios b/a aud σα at a range of timesteps for the highresolution. zero velocity dispersion halo.," The resulting axis ratio profiles can be seen in Figure \ref{fig:nice_ar}, which portrays the evolution of the axis ratios $b/a$ and $c/a$ at a range of timesteps for the high–resolution, zero velocity dispersion halo."
 The ceuter of the halo quickly becomes triaxial and remains so throughout the collapse., The center of the halo quickly becomes triaxial and remains so throughout the collapse.
 The trianialitv is fairly prolate (ο~0.725) aud appears as a central bar.," The triaxiality is fairly prolate $c \sim
0.72b$ ) and appears as a central bar."
 The outer edges of the halo have vot to collapse. and so remain spherical.," The outer edges of the halo have yet to collapse, and so remain spherical."
 As the collapse progresses. the triaxial region of the halo grows outwards. and eventually the majority of the halo settles to au equilibrium bf ratio of 0.6.," As the collapse progresses, the triaxial region of the halo grows outwards, and eventually the majority of the halo settles to an equilibrium $b/a$ ratio of $\sim0.6$."
 If the formation of the triaxial bar in Figure 2 is due to the ROI. then it should be suppressed when the orbits are less radial.," If the formation of the triaxial bar in Figure \ref{fig:nice_ar} is due to the ROI, then it should be suppressed when the orbits are less radial."
 Our simulations with larger isotropic initial velocity dispersions would then be expected to show weaker triaxialitv., Our simulations with larger isotropic initial velocity dispersions would then be expected to show weaker triaxiality.
 The particles iu these simulations have larger taugeutial velocities. which inereases the typical angular momentum of particles orbits. suppressing the formation of box orbits. and thus decreasing the streneth of the ROI," The particles in these simulations have larger tangential velocities, which increases the typical angular momentum of particles' orbits, suppressing the formation of box orbits, and thus decreasing the strength of the ROI."
 Fieure 3. shows the resulting fully evolved axis ratio profiles for a series of simulations with increasing initial raudonm velocities., Figure \ref{fig:axis_ratios} shows the resulting fully evolved axis ratio profiles for a series of simulations with increasing initial random velocities.
 The standard aud high resolution ruus with zero velocity dispersion show the central triaxial siguature of the ROL, The standard and high resolution runs with zero velocity dispersion show the central triaxial signature of the ROI.
 The le halo also shows some of this structure. but with a rounder core: this halo appears to uudersgo a weak form of the ROT.," The $1\sigma$ halo also shows some of this structure, but with a rounder core; this halo appears to undergo a weak form of the ROI."
 The 26 and 30 halos. ou the other haud. remain uearly spherical throughout uutil the eud of their evolution.," The $\sigma$ and $\sigma$ halos, on the other hand, remain nearly spherical throughout until the end of their evolution."
 Their final shapes are characterized bv rather spherical cores. with the mner axis ratios becoming progressively more spherical witli increasing initial tangential velocities.," Their final shapes are characterized by rather spherical cores, with the inner axis ratios becoming progressively more spherical with increasing initial tangential velocities."
 A preexisting isotropic velocity distribution therefore prevents the onset of selferavitating instabilities such as the ROI., A pre–existing isotropic velocity distribution therefore prevents the onset of self–gravitating instabilities such as the ROI.
 Qur results are supported by Treuti&Bertin(2006).. who find that an isotropic core acts to stabilize a radially collapsing svstem against the ROT. even when the halo is hiehlv anisotropic overall.," Our results are supported by \citet{Trenti06}, who find that an isotropic core acts to stabilize a radially collapsing system against the ROI, even when the halo is highly anisotropic overall."
 Although we do not show the full time evolution. the halos that develop a strong or moderate ROI do so quickly. establishing a prolate," Although we do not show the full time evolution, the halos that develop a strong or moderate ROI do so quickly, establishing a prolate"
Thus. the physical couditious im this cuviromment (.c.. the density. temperature aud maguctic field) are close to the range required for the σας to enit cvclotrou/svuchrotron radiation at GIIz waveleneths.,"Thus, the physical conditions in this environment (i.e., the density, temperature and magnetic field) are close to the range required for the gas to emit cyclotron/synchrotron radiation at GHz wavelengths."
 A potential problem. however. is that the dark cluster potential is shallow inside the core. which makes the state variables depend ouly weakly ou rc. unlike the steep eradieuts that are apparently required by a spectral ft with a superposition of many thermal svuchrotron couponcuts.," A potential problem, however, is that the dark cluster potential is shallow inside the core, which makes the state variables depend only weakly on $r$, unlike the steep gradients that are apparently required by a spectral fit with a superposition of many thermal synchrotron components."
 Nonetheless. the fact that the eas is not in hivdrostatie equilibrium may introduce steeper eracicuts due to its clvnamic structure. aud so definitive couclusious regarding the viability of this model to produce the GIIz spectrum of Ser A* mast be based on more detailed lvdrodvuamic simulations. which we describe next.," Nonetheless, the fact that the gas is not in hydrostatic equilibrium may introduce steeper gradients due to its dynamic structure, and so definitive conclusions regarding the viability of this model to produce the GHz spectrum of Sgr A* must be based on more detailed hydrodynamic simulations, which we describe next."
 Iu the absence of any outflow. many of Ser A*’s ridiative characteristics should be due to the deposition of the cnerey in the Galactic ceuter wind iuto the ceutral well.," In the absence of any outflow, many of Sgr A*'s radiative characteristics should be due to the deposition of the energy in the Galactic center wind into the central well."
" In the classical Boudi-Itoxle (BID) scenario (Bondi&IHovle 19113). the mass accretion rate for a iuniforii hivpersonie flow past a centralized mass Is where Ry=a2GTOtAL0,7 PEis the acerctionB radiusB and M. is. the gravitating+ mass."," In the classical Bondi-Hoyle (BH) scenario \cite{BH44}) ), the mass accretion rate for a uniform hypersonic flow past a centralized mass is where $R_A \equiv 2 G M / {v_w}^2$ is the accretion radius and $M$ is the gravitating mass."
 At the Calactic. center. for the conditions described iu the Iutroduction. we would therefore expect au accretion rate ΑμPEN10722es towith: a capture radius: Ra—0401.50.02Epe.," At the Galactic center, for the conditions described in the Introduction, we would therefore expect an accretion rate $\dot M_{BH} \sim 10^{21-22} \gms$ , with a capture radius $R_A \sim 0.01-0.02 \pc$."
 Sincem this: accretion. rate is. sub-Eddinetou. for a  one nülliou solar mass concentration. the accreting gas is mostly uuinupeded by the escapiug radiation field auc is thus esscutially iu lyvdrodvnamic frec-fall starting at Ay.," Since this accretion rate is sub-Eddington for a $\sim$ one million solar mass concentration, the accreting gas is mostly unimpeded by the escaping radiation field and is thus essentially in hydrodynamic free-fall starting at $R_A$."
 Our initial uuucerical sinulations of this process (for a poiut object). assumiue a highly simplistic wuiform flow 1991: Coker&Melia 1996)) have verified these expectations.," Our initial numerical simulations of this process (for a point object), assuming a highly simplistic uniform flow \cite{RM94}; \cite{CM96}) ) have verified these expectations."
 The Galactic center wind. however. is unlikely to bo uuiforui since many stars contribute to the mass ejection.," The Galactic center wind, however, is unlikely to be uniform since many stars contribute to the mass ejection."
 So for these calculations. we assuie that the carly-type stars euclosed. (in projection) within the Western Arc. the Northern Avia. aud the Dar produce the observed wind.," So for these calculations, we assume that the early-type stars enclosed (in projection) within the Western Arc, the Northern Arm, and the Bar produce the observed wind."
 Thus far. 25 such stars have been identified (Genzel.etal. 1996)). though the stellar wind characteristics of only 8 have been determined frou their Πο I line emission (Najarro.etal.1997: see Table 1).," Thus far, 25 such stars have been identified \cite{genz96}) ), though the stellar wind characteristics of only 8 have been determined from their He I line emission \cite{N97}; see Table 1)."
 Two of those sources. IRS 13E1 aud IRS TW. seen to dominate the mass outflow with their highwind velocity (~1000kms Ly and a mass loss rate," Two of those sources, IRS 13E1 and IRS 7W, seem to dominate the mass outflow with their highwind velocity $\sim 1000 \kms$ ) and a mass loss rate"
http://www.astrouw.edu.pl/asas/. After removing observations of declared low quality (C. D) and apparent outliers. we obtained a set of 539 /-band and 920 V-band measurements well covering the periods 1997-1999) and 2000-2009. respectively.,"http://www.astrouw.edu.pl/asas/. After removing observations of declared low quality (C, D) and apparent outliers, we obtained a set of 539 $I$ -band and 920 $V$ -band measurements well covering the periods 1997–1999 and 2000–2009, respectively."
 The robotic telescope Pi of the Sky has been designed to monitor a significant fraction of the sky with good time resolution., The robotic telescope Pi of the Sky has been designed to monitor a significant fraction of the sky with good time resolution.
 The final detector consists of two sets of 16 cameras. one camera covering a field of view of 20°x20°.," The final detector consists of two sets of 16 cameras, one camera covering a field of view of $^\circ \times
20^\circ$."
 The set of the 1101412 measurements covers the time interval of 2006-2009., The set of the 101412 measurements covers the time interval of 2006–2009.
 More details can be found at http://grb.fuw.edu.pl/. CCD photometry was done without any filter. so that the results are similar to a broad-band R colour.," More details can be found at http://grb.fuw.edu.pl/. CCD photometry was done without any filter, so that the results are similar to a broad-band $R$ colour."
 The star was also observed over eight nights in 2010 April employing the mm reflector with the classical photometer in SAAO in UBV. using a fairly conventional single channel photometer with à Hamamatsu R943-02 GaAs tube.," The star was also observed over eight nights in 2010 April employing the m reflector with the classical photometer in SAAO in $\mathit{UBV}$, using a fairly conventional single channel photometer with a Hamamatsu R943–02 GaAs tube."
 We obtained 78 triads of measurements with an inner accuracy of 9.4. 6.4. and mmmag. respectively.," We obtained 78 triads of measurements with an inner accuracy of 9.4, 6.4, and mmag, respectively."
 The very good initial estimate of the period allows us to describe the observed periodic variations by a series of phenomenological models described by à minimum number of free parameters. including the period P and the origin of epoch counting My.," The very good initial estimate of the period allows us to describe the observed periodic variations by a series of phenomenological models described by a minimum number of free parameters, including the period $P$ and the origin of epoch counting $M_0$."
 The behaviour of the light curves in V and 7 is nearly the same., The behaviour of the light curves in $V$ and $I$ is nearly the same.
 For simplicity we assumed a similar behaviour also for light curves in U and B bands. while the light curve R behaves differently.," For simplicity we assumed a similar behaviour also for light curves in $U$ and $B$ bands, while the light curve $R$ behaves differently."
 The periodic component in light variations can then be described by means of periodic functions F(Ó) and FrpeJH The functions FQ?) and Fj(éy) are the simplest normalised periodic function. that represent the observed photometric variations in detail., The periodic component in light variations can then be described by means of periodic functions $F(\vartheta)$ and $F_R(\vartheta_R)$: The functions $F(\vartheta)$ and $F_R(\vartheta_R)$ are the simplest normalised periodic function that represent the observed photometric variations in detail.
 The phase of the brightness extreme is defined to be 0.0. and the effective amplitude is defined to be 1.0.," The phase of the brightness extreme is defined to be 0.0, and the effective amplitude is defined to be 1.0."
 The functions. being the sum of three terms. are described by two dimensionless parameters Bj.B» and fi.By.," The functions, being the sum of three terms, are described by two dimensionless parameters $\beta_1,\,\beta_2$ and $\beta_3,\,\beta_4$."
 9 and ὃν are the phase function., $\vartheta$ and $\vartheta_R$ are the phase function.
 Assuming linear ephemeris. we respectively obtain: where Afe is the phase shift of the basic minimum of the function £j versus zero phase.," Assuming linear ephemeris, we respectively obtain: where $\Delta f_R$ is the phase shift of the basic minimum of the function $F_R$ versus zero phase."
 Periodic changes of magnitudes in οὐB.V.F: mj(r) and changes in R: sip(t) are given by the relations: where A is the semiamplitude of light changes common to all bands and Ax is the amplitude of an additional component of light variability being non-zero only in R.," Periodic changes of magnitudes in $U,\,B,\,V,\,I$: $m_j(t)$ and changes in $R$: $m_R(t)$ are given by the relations: where $A$ is the semiamplitude of light changes common to all bands and $A_R$ is the amplitude of an additional component of light variability being non-zero only in $R$."
 77; and 7g are mean magnitudes in the individual bands., $\overline{m_j}$ and $\overline{m_R}$ are mean magnitudes in the individual bands.
" The periodic variations of the mean value of the longitudinal component of the magnetic field intensity (B-) derived from all lines can be well approximated by the simple cosinusotd: where (B-) is the mean magnetic field intensity. A,, 1s the semiamplitude of the variations. and Afj, 8))"," The periodic variations of the mean value of the longitudinal component of the magnetic field intensity $\left<B_z\right>$ derived from all lines can be well approximated by the simple cosinusoid: where $\overline{\left<B_z\right>}$ is the mean magnetic field intensity, $A_m$ is the semiamplitude of the variations, and $\Delta f_m$ \ref{fig:longterm}) \\cite{slavek10})"
" The periodic variations of the mean value of the longitudinal component of the magnetic field intensity (B-) derived from all lines can be well approximated by the simple cosinusotd: where (B-) is the mean magnetic field intensity. A,, 1s the semiamplitude of the variations. and Afj, 8))."," The periodic variations of the mean value of the longitudinal component of the magnetic field intensity $\left<B_z\right>$ derived from all lines can be well approximated by the simple cosinusoid: where $\overline{\left<B_z\right>}$ is the mean magnetic field intensity, $A_m$ is the semiamplitude of the variations, and $\Delta f_m$ \ref{fig:longterm}) \\cite{slavek10})"
metal-enhanced wind is very powerful in reducing the metallicity of the galaxy.,metal-enhanced wind is very powerful in reducing the metallicity of the galaxy.
" In Fig. 11,,"," In Fig. \ref{Fig:mwdZmuY2},"
" models with different strengths (Amw=0.2,1,3, 10) of highly enriched winds (wyHe= 0.1) are shown."," models with different strengths $\lambda_{mw}=0, 0.2, 1, 3, 10$ ) of highly enriched winds $w_{\rm
H,He}=0.1$ ) are shown."
 We plot the evolutionary tracks of models with same bursting history as in Fig. 10., We plot the evolutionary tracks of models with same bursting history as in Fig. \ref{Fig:mwdZmuY1}.
". A stronger wind not only reduces the abundances and increases the ratio between the long recycling term elements and the short ones, but also decreases the gas fraction dramatically."," A stronger wind not only reduces the abundances and increases the ratio between the long recycling term elements and the short ones, but also decreases the gas fraction dramatically."
" It is worth to point out that although the wind efficiencies Aj, adopted here are much higher than the ones of normal wind case Aw, the gas is lost less effectively."," It is worth to point out that although the wind efficiencies $\lambda_{mw}$ adopted here are much higher than the ones of normal wind case $\lambda_{w}$, the gas is lost less effectively."
" To compare the results of normal and metal-enhanced wind models, one should assume a larger Am for the latter case."," To compare the results of normal and metal-enhanced wind models, one should assume a larger $\lambda_{mw}$ for the latter case."
" For example, Amw=10 for metal-enhanced winds is then multiplied by wyHe=0.1, therefore it is comparable with the case normal wind and λω= 1."," For example, $\lambda_{mw}=10$ for metal-enhanced winds is then multiplied by $w_{\rm H,He}=0.1$, therefore it is comparable with the case normal wind and $\lambda_{w}=1$ ."
 In the lower right panels of both Fig., In the lower right panels of both Fig.
 10 and Fig., \ref{Fig:mwdZmuY1} and Fig.
" 11we show the Y —(O/H) relations of galaxies with different degrees of enriched winds wyHe and different wind efficiencies Amw, but the same formation histories."," \ref{Fig:mwdZmuY2} we show the $Y-$ (O/H) relations of galaxies with different degrees of enriched winds $w_{\rm H,He}$ and different wind efficiencies $\lambda_{mw}$, but the same formation histories."
" By comparing with Fig. 8,"," By comparing with Fig. \ref{Fig:wdZmuY},"
 the present-day oxygen abundances are lower as we expect., the present-day oxygen abundances are lower as we expect.
" In this scenario, a very high helium abundance can be reached at low metallicity level, especially when the wind is very strong, because most of it stays inside the galaxy while heavy elements are lost."," In this scenario, a very high helium abundance can be reached at low metallicity level, especially when the wind is very strong, because most of it stays inside the galaxy while heavy elements are lost."
" The present-day Y—(O/H) relation predicted by these models do not stay on a straight line in the low metellicity region, even if the unrealistic models (very strong wind Amw=10 cases) are ruled out considering their disagreement with other observational constraints."," The present-day $Y-$ (O/H) relation predicted by these models do not stay on a straight line in the low metellicity region, even if the unrealistic models (very strong wind $\lambda_{mw}=10$ cases) are ruled out considering their disagreement with other observational constraints."
" Therefore, the observed scatter of Y—(O/H) relation may be caused by metal-enhanced winds."," Therefore, the observed scatter of $Y-$ (O/H) relation may be caused by metal-enhanced winds."
" Based on our model predictions, we suggest to fit the lower envelop of the observational data, when one derives the primordial helium Y,, because it may not be affected by the wind, hence the extrapolation to Z—0 will be more close to the real Y,"," Based on our model predictions, we suggest to fit the lower envelop of the observational data, when one derives the primordial helium $Y_p$ , because it may not be affected by the wind, hence the extrapolation to $Z=0$ will be more close to the real $Y_p$."
 Fig., Fig.
 12 is the same as Fig., \ref{Fig:mwdMsZ} is the same as Fig.
 9 but for models with metal-enhanced wind., \ref{Fig:wdMsZ} but for models with metal-enhanced wind.
" Unlike in the normal wind case, the metal-enhanced one is very effective in reducing the oxygen abundance, hence in creating the M—Z relation."," Unlike in the normal wind case, the metal-enhanced one is very effective in reducing the oxygen abundance, hence in creating the $M-Z$ relation."
" The stronger the wind efficiency, the steeper the predicted M—Z."," The stronger the wind efficiency, the steeper the predicted $M-Z$."
" Therefore, models with an increasing wind efficiency to less massive galaxies are consistent with the observations very well."," Therefore, models with an increasing wind efficiency to less massive galaxies are consistent with the observations very well."
" As a conclusion, the observed M—Z relation could be caused by metal-enhanced winds with mild strength (e.g., AmwX1)."," As a conclusion, the observed $M-Z$ relation could be caused by metal-enhanced winds with mild strength (e.g., $\lambda_{mw}\le 1$ )."
" In summary, metal-enhanced winds should take place in late-type dwarf galaxies, and they play an important role in removing gas, especially the metals, out of the galaxies."," In summary, metal-enhanced winds should take place in late-type dwarf galaxies, and they play an important role in removing gas, especially the metals, out of the galaxies."
" As we have shown in the last section, highly enriched or very strong winds will reduce the galactic oxygen abundance to a very low value during the interburst time which has not been confirmed from the observational point of view."," As we have shown in the last section, highly enriched or very strong winds will reduce the galactic oxygen abundance to a very low value during the interburst time which has not been confirmed from the observational point of view."
" Thus, in our best models wyHe=0.3 and Amw=0.8 are assumed."," Thus, in our best models $w_{\rm
H,He}=0.3$ and $\lambda_{mw}=0.8$ are assumed."
" Different numbers of bursts are examined for M;;;=10°Mo galaxies, and the best fit models are shown in Fig."," Different numbers of bursts are examined for $M_{inf}=10^9$ galaxies, and the best fit models are shown in Fig."
 13 and Fig. 14.., \ref{Fig:bestabund} and Fig. \ref{Fig:bestmuY}.
" In these models, the same SFEs (e= 0.5) and burst durations (d=0.3) are assumed, and the details of each model are listed in Table. 3.."," In these models, the same SFEs $\epsilon=0.5$ ) and burst durations (d=0.3) are assumed, and the details of each model are listed in Table. \ref{Tab:best}."
 We show the abundance ratios of different elements relative to oxygen or iron as predicted by our best models for Ming=10°Mo in Fig. 13.., We show the abundance ratios of different elements relative to oxygen or iron as predicted by our best models for $M_{inf}=10^9$ in Fig. \ref{Fig:bestabund}.
" The evolutionary tracks of these models can pass through the DLA data at early time, and cover the regions where most of BCDs have been observed."," The evolutionary tracks of these models can pass through the DLA data at early time, and cover the regions where most of BCDs have been observed."
 The evolutionary tracks of the µ—Z relation (upper panel) and Z—Y relation (lower panel) predicted by our best models are shown in Fig. 14.., The evolutionary tracks of the $\mu-Z$ relation (upper panel) and $Z-Y$ relation (lower panel) predicted by our best models are shown in Fig. \ref{Fig:bestmuY}. .
" After the wind develops, more are thebursts that a galaxy suffers, more are the oscillations in the evolutionary tracks, and thesetracks pass through most of the data, thus explaining the"," After the wind develops, more are thebursts that a galaxy suffers, more are the oscillations in the evolutionary tracks, and thesetracks pass through most of the data, thus explaining the"
Botha aud in Fig.,Both and in Fig.
 5 are caused by.ZZ. the former by the pile up of material during this pulse. and the latter when the pulse has finished.," 5 are caused by, the former by the pile up of material during this pulse, and the latter when the pulse has finished."
 They merge ~ 18.000 later. formingab.," They merge $\sim$ 18,000 later, forming."
 Shell propagates outwards. and finally disappears only slightly moclilvine the surrounding eas.," Shell propagates outwards, and finally disappears only slightly modifying the surrounding gas."
 Shells and have (the same formation process as shells audb mentioned previously. but in this case their origin is associated with///.," Shells and have the same formation process as shells and mentioned previously, but in this case their origin is associated with."
 Again. because is being decelerated. catches it up in ~3xLO! and hence appears.," Again, because is being decelerated, catches it up in $\sim 3\times10^4$ and hence appears."
 In Fig., In Fig.
 6 we have shown the gas density and velocity evolution al times which correspond in Fie., 6 we have shown the gas density and velocity evolution at times which correspond in Fig.
 1 to 4.500 and 3.200 before and after the end of respectively and 4.500. 16.500 and 71.300 after the beginning ofV.," 1 to 4,500 and 3,200 before and after the end of respectively and 4,500, 16,500 and 71,300 after the beginning of."
 At the top of Fig., At the top of Fig.
 6 we have marked (he density peak which is formed byZV., 6 we have marked the density peak which is formed by.
 A strong shock develops in the density valley formed between ande. leading to the Formation ofx.," A strong shock develops in the density valley formed between and, leading to the formation of."
 When pulsefinishes. shellf appears.," When pulse finishes, shell appears."
 With increasing time. all the density. peaks(ed.e. £f) disappear because {μον [eedx.," With increasing time, all the density peaks, ) disappear because they feed."
 Only remains as an identifiable shell until the end of the AGB., Only remains as an identifiable shell until the end of the AGB.
 PulseV is responsible for the increase labeled as which completely disappears 4.400 after when the density. falls as a power law of the type r7 Jaw.," Pulse is responsible for the increase labeled as which completely disappears 4,400 after when the density falls as a power law of the type $r^{-2}$ law."
 This is what is expected for a constant mass-loss rate al a steady. wind velocity., This is what is expected for a constant mass-loss rate at a steady wind velocity.
 The disruption of the density structure seen in the bottom plot of Fig., The disruption of the density structure seen in the bottom plot of Fig.
 6 is caused bv the velocity decrease produced during the middle of pulseV., 6 is caused by the velocity decrease produced during the middle of pulse.
 Of all the cases considered this model has the longest evolution during the AGB., Of all the cases considered this model has the longest evolution during the AGB.
 The wind parameters undergo (11ος main mocdulations in mass-loss rate and velocity. which occur during the last 300.000 (see Fie.," The wind parameters undergo three main modulations in mass-loss rate and velocity, which occur during the last 300,000 (see Fig."
 1)., 1).
 The evolution of the gas is displaved at five snapshots during the TP-AGB in Fig., The evolution of the gas is displayed at five snapshots during the TP-AGB in Fig.
 7., 7.
 The times selected have a relation to the stellar evolution as follows: the top plot corresponds to 6.500 alter the end ofZ£. the second and third plots have been selected at 39.000 and 8.000 after the beginine ancl end of£77 respectively. ancl the fourth and fifth plots at 45.500 and 12.300 alter the beeining ofJV.," The times selected have a relation to the stellar evolution as follows: the top plot corresponds to 6,500 after the end of, the second and third plots have been selected at 39,000 and 8,000 after the begining and end of respectively, and the fourth and fifth plots at 45,500 and 12,300 after the begining of."
 The sinall modulations in velocity experienced by (he wind before pulse7/ do not have anv effect on the gas structure., The small modulations in velocity experienced by the wind before pulse do not have any effect on the gas structure.
 The first significant effect arises as a consequence ofLf when a high density shell is formed., The first significant effect arises as a consequence of when a high density shell is formed.
 Due to its high thermal pressure (his shell expands splitting the density aud velocity into (wo peaks., Due to its high thermal pressure this shell expands splitting the density and velocity into two peaks.
 This peculiar two peak structure lasts lor around 22.000vr. disappearing as the eas propagates outwards. resulting in (he formation of (see top plot in Fig.," This peculiar two peak structure lasts for around 22,000, disappearing as the gas propagates outwards, resulting in the formation of (see top plot in Fig."
 7)., 7).
 Shells and are formed by the gradual end of// first. both the wind velocity and the mass-loss rate are reduced by factors of 3 and 5 respectively is lormed). and second. the mass-Ioss rate drops again a [actor of 1.5. meanwhile the wind velocity grows inf/f and is formed.," Shells and are formed by the gradual end of: first, both the wind velocity and the mass-loss rate are reduced by factors of 3 and 5 respectively is formed), and second, the mass-loss rate drops again a factor of 1.5, meanwhile the wind velocity grows in and is formed."
 Shell survives for only à short time (3.000 vr). i( is accelerated," Shell survives for only a short time (8,000 yr), it is accelerated"
agreement between moclel predictions ancl observations concerning all the properties of the observed abundance profiles (absolute values. gradient. scatter) for O. S. Ne and Ar.,"agreement between model predictions and observations concerning all the properties of the observed abundance profiles (absolute values, gradient, scatter) for O, S, Ne and Ar."
 The model suggests that abundance gradients are sleeper in (he earlier epoch., The model suggests that abundance gradients are steeper in the earlier epoch.
 However. the large scatter in the adopted data does not allow one to conclude on the temporal variation of (he gradients.," However, the large scatter in the adopted data does not allow one to conclude on the temporal variation of the gradients."
 Nevertheless PNs sulfer from laree uncertainties concerning (heir progenitor's masses and lifetimes as well as their distances from Galactic center., Nevertheless PNs suffer from large uncertainties concerning their progenitor's masses and lifetimes as well as their distances from Galactic center.
 On the other hand. open clusters(OCs) have long been used (ο trace the structure ancl evolution of the Galactic disk (Friel1995).," On the other hand, open clusters(OCs) have long been used to trace the structure and evolution of the Galactic disk \citep{fri95}."
. Since open clusters could be relatively accurately dated ancl we can see them to large distance. their [Fe/1I] values serve an excellent tracer to the abundance gradient along the Galactic disk as well as many other important disk properties. such as Age-Metallicity Relation(AMIR). abundance gradient evolution. disk age and so on (Carraroetal.1993).," Since open clusters could be relatively accurately dated and we can see them to large distance, their [Fe/H] values serve an excellent tracer to the abundance gradient along the Galactic disk as well as many other important disk properties, such as Age-Metallicity Relation(AMR), abundance gradient evolution, disk age and so on \citep{car98}."
. At this point. one might ask whether the field disk populations are also able to (trace the disk evolution.," At this point, one might ask whether the field disk populations are also able to trace the disk evolution."
 Indeed. the extensive studies by Exbvrdssonetal.(1993)... and recently by Chenetal.(2000).. who concentrate on disk EF. G stars. show an overall radial gradient. (hat is nearly independent of age.," Indeed, the extensive studies by \citet{edv93}, and recently by \citet{che00}, who concentrate on disk F, G stars, show an overall radial gradient that is nearly independent of age."
 Those results are based on stars mainly restricted in (he solar neighborhood., Those results are based on stars mainly restricted in the solar neighborhood.
 A more detailed analvsis for the disk iron gradient was given by on the basis of 1302 field star with high resolution proper motion and parallax data from Iipparcos satellite., A more detailed analysis for the disk iron gradient was given by \citet{cui00} on the basis of 1302 field star with high resolution proper motion and parallax data from Hipparcos satellite.
 Thev have derived an radial iron gradient of —0.057 wwithin galactocentric distance from 8.5 kpe to 17 kpe., They have derived an radial iron gradient of $-$ 0.057 within galactocentric distance from 8.5 kpc to 17 kpc.
 However. it is still diffieult to reveal anv pronounced gradient evolution from those results.," However, it is still difficult to reveal any pronounced gradient evolution from those results."
 Moreover. results from those studies are strongly affected by selection effects and rely on the techniques lor determining individual stellar distances (which are heavily dependent on the adopted Galaxy. potential model) that are much less reliable than those used to obtain cluster distances.," Moreover, results from those studies are strongly affected by selection effects and rely on the techniques for determining individual stellar distances (which are heavily dependent on the adopted Galaxy potential model) that are much less reliable than those used to obtain cluster distances."
 In a recent work. have modelled the effects of the orbital diffusion of stars and. elusters on the Galactic abundance gradient.," In a recent work, \citet{cor01} have modelled the effects of the orbital diffusion of stars and clusters on the Galactic abundance gradient."
 The general conclusion is that the effect of diffusion makes a gradient shallower over lime. and the cluster population offers a more viable means for lincdine detailed structure within the recent Galactic abundance gradient.," The general conclusion is that the effect of diffusion makes a gradient shallower over time, and the cluster population offers a more viable means for finding detailed structure within the recent Galactic abundance gradient."
 llere. we also pointed out that our recent treatinent on deriving the abundance gradient from open clusters in Houetal.(2002) is in [act not proper.," Here, we also pointed out that our recent treatment on deriving the abundance gradient from open clusters in \citet{hou02} is in fact not proper."
 In that paper. we have simply taken low Catalogs from literatures (Carraroetal.1998;TwarogAshmanandAnthonv-Dwarog1997:Piattietal.1995:Friel 1995).. and merge them just by making cross checking for the common clusters. without examining individually to see if there are important difference among clusters in the different catalogs (Twarog2002).," In that paper, we have simply taken four Catalogs from literatures \citep{car98, twa97, pia95, fri95}, and merge them just by making cross checking for the common clusters, without examining individually to see if there are important difference among clusters in the different catalogs \citep{twa02}."
. We reler this paper to act as a substitution (ο our previous one., We refer this paper to act as a substitution to our previous one.
 In this paper. we compiled a set of new open cluster catalogues.," In this paper, we compiled a set of new open cluster catalogues."
 The catalogue was, The catalogue was
adequate fit to most of the velocities.,adequate fit to most of the velocities.
 There is a significant departure from this sinusoidal behaviour for the few points prior to the eclipse., There is a significant departure from this sinusoidal behaviour for the few points prior to the eclipse.
" This is in line with a disc origin for these absorption features since at these phases the line-of-sight will be through the bright spot, and so our measurements will be distorted by the complex dynamics of the disc-stream interaction region."," This is in line with a disc origin for these absorption features since at these phases the line-of-sight will be through the bright spot, and so our measurements will be distorted by the complex dynamics of the disc-stream interaction region."
" As well as the velocity, we note that the lines seem to vary in strength with phase: Figure 6 shows them to be fainter (and also slightly broader) in the few spectra leading up to phase 0.5 (N= 8) and phase 0 19), but much stronger in the intermediate phases."," As well as the velocity, we note that the lines seem to vary in strength with phase: Figure \ref{fig:uvb_trail} shows them to be fainter (and also slightly broader) in the few spectra leading up to phase $0.5$ $N=8$ ) and phase $0$ $N=19$ ), but much stronger in the intermediate phases."
" After removing the orbital velocity shifts and averaging the spectra, we fitted a subset of the iron lines and determined their full-width half maximum (FWHM) to be 85.5+2.9 km/s. For this subset we excluded lines which were potentially blended."," After removing the orbital velocity shifts and averaging the spectra, we fitted a subset of the iron lines and determined their full-width half maximum (FWHM) to be $85.5 \pm 2.9$ km/s. For this subset we excluded lines which were potentially blended."
 We find no significant difference, We find no significant difference
to the diffuse extragalactic gamma-ray background (EGB). which was also measured by EGRET fora subtle issue of Galactic foregroundsubtraction).,"to the diffuse extragalactic gamma-ray background (EGB), which was also measured by EGRET for a subtle issue of Galactic foreground."
 This paper ts organized as follows., This paper is organized as follows.
 In 2.. we summarize the predictions of SSC model for the prompt 2.1) and afterglow 2.2)) phases.," In \ref{sec:IC}, we summarize the predictions of SSC model for the prompt \ref{sub:prompt}) ) and afterglow \ref{sub:afterglow}) ) phases."
 Section 3. 1s devoted for analysis of the GRB fluence data by EGRET. from which distributions of fluence in the GeV band are derived.," Section \ref{sec:Constraint on
high-energy emission with EGRET} is devoted for analysis of the GRB fluence data by EGRET, from which distributions of fluence in the GeV band are derived."
 We then use these distributions to argue prospects for GRBdetection withGLAST in 4.. and implications for EGB from GRB emissions in 5..," We then use these distributions to argue prospects for GRBdetection with in \ref{sec:GLAST}, and implications for EGB from GRB emissions in \ref{sec:EGB}."
 In 6.. we give a summary of the present paper.," In \ref{sec:conclusions}, we give a summary of the present paper."
 If the prompt and/or afterglow emission is due to synchrotron radiation from relativistic electrons (with Lorentz factor το). then there must be an accompanying IC component from the same electrons scattering off the synchrotron photons.," If the prompt and/or afterglow emission is due to synchrotron radiation from relativistic electrons (with Lorentz factor $\gamma_e$ ), then there must be an accompanying IC component from the same electrons scattering off the synchrotron photons."
 The spectral shape of the IC emission is almost the same as the synchrotron radiation (shifted by D. and is expected to fall around the GeVrange during both the prompt and afterglow phases.," The spectral shape of the IC emission is almost the same as the synchrotron radiation (shifted by $\gamma_e^2$ ), and is expected to fall around the GeVrange during both the prompt and afterglow phases."
" For e,>eg. and assuming that there is no ""Kleimn-Nishina. suppression"" and that the emitting electrons are fast cooling. the IC fluence is related to the synchrotron fluence simply through Ficz(ενερ)!YEa."," For $\epsilon_e > \epsilon_B$, and assuming that there is no “Klein-Nishina suppression” and that the emitting electrons are fast cooling, the IC fluence is related to the synchrotron fluence simply through $F_{\rm IC} \approx (\epsilon_e / \epsilon_B)^{1/2}
F_{\rm syn}$."
 Thus. assuming that the microphysics do not vary much from burst to burst. it is natural to assume proportionality between the synchrotron MeV fluence (observed by BATSE) and the GeV synchrotron plus IC fluence (observed by EGRET and in the future by GLAST): where ;jj44 and (jc are coefficients for the proportionality due to synchrotron and IC processes.," Thus, assuming that the microphysics do not vary much from burst to burst, it is natural to assume proportionality between the synchrotron MeV fluence (observed by BATSE) and the GeV synchrotron plus IC fluence (observed by EGRET and in the future by ): where $\eta_{\rm syn}$ and $\eta_{\rm IC}$ are coefficients for the proportionality due to synchrotron and IC processes."
 Note that the synchrotron fluence in the GeV range can be extrapolated relatively easily. if we assume that the spectrum extends up to such high energies.," Note that the synchrotron fluence in the GeV range can be extrapolated relatively easily, if we assume that the spectrum extends up to such high energies."
 Thus. we here focus on theoretical evaluation of the [C component.," Thus, we here focus on theoretical evaluation of the IC component."
" At first approximation. the coefficient jc is roughly (e.ερ)! from considerations above. and thus we define where for the prompt emission Z4 Fyey while for the afterglow F., is the afterglow fluence within the radio to X-ray energy bands."," At first approximation, the coefficient $\eta_{\rm IC}$ is roughly $(\epsilon_e / \epsilon_B)^{1/2}$ from considerations above, and thus we define where for the prompt emission $F_{\rm syn} \approx F_{\rm MeV}$ while for the afterglow $F_{\rm syn}$ is the afterglow fluence within the radio to X-ray energy bands."
" Correction factors {gy and ἐν represent the effect of Klein-Nishina suppression and detector energy window, respectively. which are given below."," Correction factors $\xi_{\rm
KN}$ and $\xi_w$ represent the effect of Klein-Nishina suppression and detector energy window, respectively, which are given below."
" We define typical frequencies for both synchrotron (444) and IC G4c) as the frequencies where most of the energies are radiated in case that the Klein-Nishina cross section does not play an important role: re.. where 7f,, for each component is peaked in this case."," We define typical frequencies for both synchrotron $\nu_{\rm syn}$ ) and IC $\nu_{\rm IC}$ ) as the frequencies where most of the energies are radiated in case that the Klein-Nishina cross section does not play an important role; i.e., where $\nu
f_\nu$ for each component is peaked in this case."
 From relativistic kinematics. these two typical frequencies are related through where σι Is a characteristic Lorentz factor of the electrons that dominate the synchrotron power1979): this is true in the fast cooling regime. which is the case in the most of our discussions).," From relativistic kinematics, these two typical frequencies are related through where $\gamma_m$ is a characteristic Lorentz factor of the electrons that dominate the synchrotron power; this is true in the fast cooling regime, which is the case in the most of our discussions."
. The Klein-Nishina effect is relevant if a photon energy in the electron rest frame exceeds the electron rest mass energy. and this condition is formulated as where Τρ is the bulk Lorentz factor of the ejecta. which is on the order of 100 in the prompt phase of GRBs and their early afterglows.," The Klein-Nishina effect is relevant if a photon energy in the electron rest frame exceeds the electron rest mass energy, and this condition is formulated as where $\Gamma_b$ is the bulk Lorentz factor of the ejecta, which is on the order of 100 in the prompt phase of GRBs and their early afterglows."
 Upscattering synchrotron photons to energies above {ην is highly suppressed. which results in IC cutoff at gw.," Upscattering synchrotron photons to energies above $h \nu_{\rm KN}$ is highly suppressed, which results in IC cutoff at $\nu_{\rm KN}$."
 Besides producing a spectral cutoff. the Klein-Nishina effect also modifies the way electrons cool. which is relevant for the GeV emission and is also included in £gw.," Besides producing a spectral cutoff, the Klein-Nishina effect also modifies the way electrons cool, which is relevant for the GeV emission and is also included in $\xi_{\rm KN}$."
 Electrons with energies above Klem-Nishina threshold (for a given seed-photon energy) can lose their energies only through synchrotron radiation. while the lower-energy ones can cool through both processes.," Electrons with energies above Klein-Nishina threshold (for a given seed-photon energy) can lose their energies only through synchrotron radiation, while the lower-energy ones can cool through both processes."
 Such an effect has been studied in the case where the seed photons for IC scattering are provided by an external sourcestherein)., Such an effect has been studied in the case where the seed photons for IC scattering are provided by an external sources.
. However. in the case of SSC mechanism. since the seed photons are emitted from synchrotron process due to the same electron population. we should properly take into account feedback.," However, in the case of SSC mechanism, since the seed photons are emitted from synchrotron process due to the same electron population, we should properly take into account feedback."
 Giving full details on this is beyond the scope of the present paper. but some results are summarized briefly in Appendix A (see also 2003)).," Giving full details on this is beyond the scope of the present paper, but some results are summarized briefly in Appendix \ref{app:KN}
 (see also )."
 Here we only show the approximate analytic form of £y: where xy 15 theLorentz factor of electrons for which photons at 72r4 are in the Klein-Nishina regime., Here we only show the approximate analytic form of $\xi_{\rm KN}$: where $\gamma_{\rm KN}$ is theLorentz factor of electrons for which photons at $\nu \gtrsim \nu_{\rm syn}$ are in the Klein-Nishina regime.
" The energy of an observed photon with frequency » as measured in the rest frame of an electron with Lorentz factor  is zzΙΤ where the 1/T, factor converts the photon energy from the observer frame to theplasma rest frame and the  factor converts it to the electron rest frame.", The energy of an observed photon with frequency $\nu$ as measured in the rest frame of an electron with Lorentz factor $\gamma$ is $\approx \gamma h\nu/\Gamma_b$ where the $1/\Gamma_b$ factor converts the photon energy from the observer frame to theplasma rest frame and the $\gamma$ factor converts it to the electron rest frame.
 Since such a photon is in the Klein-Nishina regime of an electron with Lorentz factor ~ once its energy in the electron rest frame is larger than Dc. we obtain: This Klein-Nishina feedbackeffect modifies the spectrum shape of both synchrotron and IC emissions (in addition to the Klein-Nishina cutoff for IC)., Since such a photon is in the Klein-Nishina regime of an electron with Lorentz factor $\gamma$ once its energy in the electron rest frame is larger than $m_e c^2$ we obtain: This Klein-Nishina feedbackeffect modifies the spectrum shape of both synchrotron and IC emissions (in addition to the Klein-Nishina cutoff for IC).
 We note that equation (5) provides a solution that agrees within a factor of —2 with the one obtained by numerically solving equation (AT))., We note that equation \ref{eq:eta KN}) ) provides a solution that agrees within a factor of $\sim$ 2 with the one obtained by numerically solving equation \ref{eq:Y}) ).
 This precision ts sufficiently good for our purpose. especially because it is well within the uncertainty ranges of other parameters.," This precision is sufficiently good for our purpose, especially because it is well within the uncertainty ranges of other parameters."
 By &..we take into account the fraction of the IC fluence that falls into the GeV detector energy bands.," By $\xi_w$,we take into account the fraction of the IC fluence that falls into the GeV detector energy bands."
" EGRET window is between /4,;=30 MeV and /n4,,,=30 GeV while GLAST-LAT window is between /n4,.,=20 MeV and /i,.,,=300 GeV. We here assume that the frequency where most of the IC energy is released. Meja—min[7jc.iiw]. 1s always larger than lower limit of the frequency band. με. as expected for both EGRET and GLAST. and thus consider the cases in which Me Is within or above the detector frequency band."," EGRET window is between $h
\nu_{w,l}=30$ MeV and $h \nu_{w,u}=30$ GeV while }-LAT window is between $h \nu_{w,l}=20$ MeV and $h \nu_{w,u}=300$ GeV. We here assume that the frequency where most of the IC energy is released, $\nu_{\rm IC, peak} \equiv \min[\nu_{\rm IC}, \nu_{\rm
KN}]$, is always larger than lower limit of the frequency band, $\nu_{w,l}$, as expected for both EGRET and , and thus consider the cases in which $\nu_{\rm IC,peak}$ is within or above the detector frequency band."
" In the formerpea, case where 14,4<<11cpoatX 14. We have &,7 I."," In the former case where $\nu_{w,l} < \nu_{\rm IC, peak} < \nu_{w,u}$ , we have $\xi_{w} \approx 1$ ."
" On the other hand. if mepeak> Meu then most of the energy comes from the upper frequency limit 7... and we have Gy©(u/MEpeak y. where a, is the photon spectral"," On the other hand, if $\nu_{\rm IC,
peak} > \nu_{w,u}$ then most of the energy comes from the upper frequency limit $\nu_{w,u}$ , and we have $\xi_{w} \approx (\nu_{w,u}
/ \nu_{\rm IC, peak})^{2 - \alpha_1}$ , where $\alpha_1$ is the photon spectral"
reality. we expect Z'(0) to be arger because of radiative heating bv the X-rays from the post-shock region above and 72 to be less than that cletermined bv the free-[all velocity at the shock surface (sco Cropper 11909).,"reality, we expect $T(0)$ to be larger because of radiative heating by the X-rays from the post-shock region above and $T_{\rm s}$ to be less than that determined by the free-fall velocity at the shock surface (see Cropper 1999)."
 Therefore. we have underestimated τι here. and he value that we obtained can be considerec as a lower limit to my.," Therefore, we have underestimated $\tau_{\rm m}$ here, and the value that we obtained can be considered as a lower limit to $\tau_{\rm m}$."
" Nevertheless. under this simple approximation. we have found that the assumption τι<Ty, is acceptable. provided that Ziyp is not. very significantly larger than 103 Ix. Moreover. the assumption that zu,«1 is in general valiL."," Nevertheless, under this simple approximation, we have found that the assumption $\tau_{\rm h} < \tau_{\rm m}$ is acceptable, provided that $T_{\rm eff}$ is not very significantly larger than $10^4$ K. Moreover, the assumption that $\tau_{\rm m} \ll 1$ is in general valid."
 Vhis is shown in Figure 6., This is shown in Figure 6.
" Wuich values of 7, can »* reached. is therefore unclear pending furtjer investigation. but the values indicated. [rom Figure 6 are in the range 10 500 10.7. significantly lower than those used in Figure 5."," Which values of $\tau_{\rm m}$ can be reached is therefore unclear pending further investigation, but the values indicated from Figure 6 are in the range $10^{-4}$ to $10^{-5}$, significantly lower than those used in Figure 5."
 Ehe dillerence between spectra from. models with these smaller values of τιν and those from the Aizu (1973) formulation is less than in the range of 0.05 eV to 100 keV. so that the observational elfects of the leaky base will benegligixe.," The difference between spectra from models with these smaller values of $\tau_{\rm m}$ and those from the Aizu (1973) formulation is less than in the range of 0.05 eV to 100 keV, so that the observational effects of the leaky base will benegligible."
 Phus i£ we can indeed rely on small values of Tu. Current cold wall models such as those mentioned in Section 1 can continue to be emploved for X-ray spectral fitting.," Thus if we can indeed rely on small values of $\tau_{\rm m}$, current cold wall models such as those mentioned in Section 1 can continue to be employed for X-ray spectral fitting."
 The advantage of our treatment is that it avoids the infinities at the base of the flow inherent in all previous formulations in whic ithe stationary-wall boundary condition base is assumed AAizu 73: Chevalier Imanmura: Wu 1994: Wu 11994)., The advantage of our treatment is that it avoids the infinities at the base of the flow inherent in all previous formulations in which the stationary-wall boundary condition base is assumed Aizu 1973; Chevalier Imamura; Wu 1994; Wu 1994).
 This is a step in the direction of more realistic formulations which allow the hvdrodynamie accretion Dow to match the ivedrostatic white-chwarl atmosphere., This is a step in the direction of more realistic formulations which allow the hydrodynamic accretion flow to match the hydrostatic white-dwarf atmosphere.
 As noted above the predicted spectral changes are small for small my: however apractical advantage is that it eliminates the numerical errors resulting from finite sampling ofquantities which tend to infinity., As noted above the predicted spectral changes are small for small $\tau_{\rm m}$; however a advantage is that it eliminates the numerical errors resulting from finite sampling of quantities which tend to infinity.
 In the oevious treatments the fineness of the sampling of the base of the post-shock region in the numerical integrating scheme has an elfect on the spectral slope: finer sampling will encounter a higher value of the density (which is tending to infinity) at the ose. so that softer spectra will be generated.," In the previous treatments the fineness of the sampling of the base of the post-shock region in the numerical integrating scheme has an effect on the spectral slope: finer sampling will encounter a higher value of the density (which is tending to infinity) at the base, so that softer spectra will be generated."
 In our formulation. as the singularity at the base is removed. the elfect of the sampling is insignificant even when small values of τι are assumed.," In our formulation, as the singularity at the base is removed, the effect of the sampling is insignificant even when small values of $\tau_{\rm m}$ are assumed."
 This is consistent with the fact that emission from the »»undary laver should be finite. as deduced from equation (45).," This is consistent with the fact that emission from the boundary layer should be finite, as deduced from equation (45)."
 We thank Mark Wardle for comments and for carefully reading the manuscript., We thank Mark Wardle for comments and for carefully reading the manuscript.
 We also thank the referee for providing helpful, We also thank the referee for providing helpful
The parameters for the spectral lines were obtained from VALD (?)..,The parameters for the spectral lines were obtained from VALD \citep{vald2}.
 To optimize the fit to the observations. the e logf- of some lines was changed.," To optimize the fit to the observations, the $\log gf$ -value of some lines was changed."
 The spectral parameters for the most important lines are given in Table 4.., The spectral parameters for the most important lines are given in Table \ref{speparam}.
 The full line synthesis included 151 lines., The full line synthesis included 151 lines.
 The need to adjust the spectral parameters could indicate that the element abundances differ from those of the standard MARCS model., The need to adjust the spectral parameters could indicate that the element abundances differ from those of the standard MARCS model.
 It should be emphasized that these kinds of corrections are necessary in Doppler imaging to reduce the systematic errors caused by discrepancies between the synthetic spectra and observations., It should be emphasized that these kinds of corrections are necessary in Doppler imaging to reduce the systematic errors caused by discrepancies between the synthetic spectra and observations.
 A table of line profiles was calculated using MARCS plane-parallel atmosphere models with Των ranging from 3200 K to 6000 K and the stellar parameters listed in Table 3.., A table of line profiles was calculated using MARCS plane-parallel atmosphere models with $T_\mathrm{eff}$ ranging from 3200 K to 6000 K and the stellar parameters listed in Table \ref{sparam}.
 The local line profiles table and the observations were used as input to the inversion code., The local line profiles table and the observations were used as input to the inversion code.
 The inversion was based on Tikhonov regularization., The inversion was based on Tikhonov regularization.
 A regularization parameter of A=I-107? and a grid of 80 x 40 surface elements were used., A regularization parameter of $\Lambda = 1 \cdot 10^{-9}$ and a grid of 80 x 40 surface elements were used.
 The aim was to achieve a difference between the model and the observations of about the same level as the observational noise. 1.e. the inverse of the S/N ratio.," The aim was to achieve a difference between the model and the observations of about the same level as the observational noise, i.e. the inverse of the $S/N$ ratio."
 This level of convergence was reached after 40 iterations., This level of convergence was reached after 40 iterations.
 The deviations are listed in Table |.., The deviations are listed in Table \ref{obssum}.
 In most cases the deviation was significantly larger than the inverse of the S /N-value., In most cases the deviation was significantly larger than the inverse of the $S/N$ -value.
 This may be caused by systematic errors. e.g. slight shifts in the continuum level or modelling errors.," This may be caused by systematic errors, e.g. slight shifts in the continuum level or modelling errors."
 In some seasons. there may also be ehanges in the spot configuration happening on a timescale shorter than a week.," In some seasons, there may also be changes in the spot configuration happening on a timescale shorter than a week."
 The resulting maps are displayed in Fig., The resulting maps are displayed in Fig.
 | and comparisons between the modelled and observed spectra are shown in Fig. 2.., \ref{dimaps} and comparisons between the modelled and observed spectra are shown in Fig. \ref{dispectra}.
 To make a more reliable comparison with our earlier Doppler images of 1994— 2002 (?). we recalculated the images from 2002 using the new stellar parameters and model atmospheres.," To make a more reliable comparison with our earlier Doppler images of 1994 – 2002 \citep{lindborg2011}, we recalculated the images from 2002 using the new stellar parameters and model atmospheres."
 Since the earlier observations did not include the same wavelength regions as the present ones. we used the same wavelength regions as in the previous study.," Since the earlier observations did not include the same wavelength regions as the present ones, we used the same wavelength regions as in the previous study."
 The resulting maps are shown in the two uppermost left panels of Fig. 1., The resulting maps are shown in the two uppermost left panels of Fig. \ref{dimaps}. .
spectrum in the optical. with the two components producing half of the [lux cach around the Z-band.,"spectrum in the optical, with the two components producing half of the flux each around the $I$ -band."
 We were allocated ΝΕΤ imaging polarimetry with FORSL (optical) for these observations (ESO Programme ID. 076.D-0497)., We were allocated VLT imaging polarimetry with FORS1 (optical) for these observations (ESO Programme ID 076.D-0497).
 The 2005 outburst of GRO J165540 was studied at X-ray. optical and radio wavelengths. (c.g.Shaposhnikovctal.2007).," The 2005 outburst of GRO J1655–40 was studied at X-ray, optical and radio wavelengths \citep[e.g.][]{shapet07}."
. On the outburst decline the source entered the [owπαντα state (Llomanetal.2005b) so we obtained ΔΙΑ imaging polarimetry with a Directors Discretionary Time (DDT) proposal with ISAAC (E6O Programme LD 275.D-5062)., On the outburst decline the source entered the low/hard state \citep{homaet05b} so we obtained NIR imaging polarimetry with a Director's Discretionary Time (DDT) proposal with ISAAC (ESO Programme ID 275.D-5062).
 In addition to outbursting sources. we selected a number of quiescent ancl persistent sources for NIIT polarimetry.," In addition to outbursting sources, we selected a number of quiescent and persistent sources for NIR polarimetry."
 Target selection. of quiescent systems was based on known NIRo magnitudes and the companion star contribution. where measured.," Target selection of quiescent systems was based on known NIR magnitudes and the companion star contribution, where measured."
 For example. the companion in the A0620-00 system is known to contribute ~754 of the IH -band quiescent lux (Shahbaz.Bandyopadhyay&Charles 1999).," For example, the companion in the A0620-00 system is known to contribute $\sim 75$ of the $H$ -band quiescent flux \citep*{shahet99}."
.. Even if a polarised jet. provides the remaining (as opposed to the disc). the level of LP would be dampened wea factor of 4 due to the unpolarised light (rom the star.," Even if a polarised jet provides the remaining (as opposed to the disc), the level of LP would be dampened by a factor of 4 due to the unpolarised light from the star."
 We found that most NIR-xight. BIINBs are dominated by heir companion stars (c.g.V404CryerCasaresetal.1993) but there are some good candidates (oe. GRS 1915|105: ATE 115|480)., We found that most NIR-bright BHXBs are dominated by their companion stars \citep[e.g. V404 Cyg;][]{casaet93} but there are some good candidates (e.g. GRS 1915+105; XTE J1118+480).
 Three SABs are also good. candidates o identify a NIR jet component., Three NSXBs are also good candidates to identify a NIR jet component.
 Optically thin svnchrotron emission appears to join o the thermal disc spectrum in he NIR in the NSXD 4U 0614|09 (Migliarictal.2006a)., Optically thin synchrotron emission appears to join to the thermal disc spectrum in the NIR in the NSXB 4U 0614+09 \citep{miglet06}.
. Sco X lis the brightest (J= 11.9) of the Z-sources. a class of neutron star whose jet spectrum may significantly contribute to the NER.," Sco X–1 is the brightest $J = 11.9$ ) of the Z-sources, a class of neutron star whose jet spectrum may significantly contribute to the NIR."
 Aql X. Lis very active (it has about one outburst per vear) and has a NIR spectrum in outburst consistent with a jet origin (Russelletal.2007)., Aql X–1 is very active (it has about one outburst per year) and has a NIR spectrum in outburst consistent with a jet origin \citep{russet07}.
. In Table 1 we list. all observations usec in this work., In Table 1 we list all observations used in this work.
 Lor all three instruments (on two telescopes) imaging polarimetry is obtained using a Wollaston prism inserted in the optical path., For all three instruments (on two telescopes) imaging polarimetry is obtained using a Wollaston prism inserted in the optical path.
 The light) is split. into simultaneous. perpendicularly polarised ordinary ancl extra- beams (o- and e-beams).," The light is split into simultaneous, perpendicularly polarised ordinary and extra-ordinary beams (o- and e-beams)."
 The waveplate was rotated to four angles: 0 .22.5. 45 and 67.5%: this achieves a higher accuracy of LP and PA than from just two angles O and 45°.," The waveplate was rotated to four angles: $^\circ$, $^\circ$, $^\circ$ and $^\circ$; this achieves a higher accuracy of LP and PA than from just two angles $^\circ$ and $^\circ$ ."
 The Dux of the source was then measured in the o- and e-beams for each prism angle GF. Po. F5S. dso. ete.)," The flux of the source was then measured in the o- and e-beams for each prism angle $F_0^o$, $F_0^e$, $F_{22}^o$, $F_{22}^e$, etc.)"
 using aperture photometry in LRA and LP and PA were calculated: using the ratio method thus: Flax calibration was possible for most of our targets., using aperture photometry in IRAF and LP and PA were calculated using the `ratio' method thus: Flux calibration was possible for most of our targets.
 The total flux of each source was estimated. from the sum of the o- and e-beam IIuxes at each angle., The total flux of each source was estimated from the sum of the o- and e-beam fluxes at each angle.
 Magnitudes were converted. to cle-redeened Dux.densities by accounting for interstellar extinction., Magnitudes were converted to de-reddened fluxdensities by accounting for interstellar extinction.
 We adopted the method of, We adopted the method of
time (Zhengetal. 1999)).,time \cite{Zheng99}) ).
 The filter used for the observations reported here is the BATC filter with central wavelength of and bandwidth of I80À (Fanetal.1996. and Yanctal.2000)).," The filter used for the observations reported here is the BATC filter with central wavelength of and bandwidth of $\rm
480{\AA}$ \cite{Fan96} and \cite{Yan00}) )."
 This filter is a eood compromise between eotting as far to the red as possible. aud avoiding many bright sky cussion lines with the broadest possible filter.," This filter is a good compromise between getting as far to the red as possible, and avoiding many bright sky emission lines with the broadest possible filter."
 Indeed. all of the BATC filters are designed to avoid contamination by emission lines from the wieht-sky (cf.," Indeed, all of the BATC filters are designed to avoid contamination by emission lines from the night-sky (cf."
 Fan ot al., Fan et al.
 1996)., 1996).
 We obtained 190 images of NGC 1565 frou 1995 to 1997. of which we deemed 150 images taken in 22 runs (Table 1) as suitable for analysis.," We obtained 190 images of NGC 4565 from 1995 to 1997, of which we deemed 150 images taken in 22 runs (Table 1) as suitable for analysis."
 Most of these images are of exposure times 900 to 1200 sec., Most of these images are of exposure times 900 to 1200 sec.
 In a given uieht all exposures were dithered raucdomly at a level of ~10 pixels to facilitate removal of cosmic rays and CCD defects during the data reduction., In a given night all exposures were dithered randomly at a level of $\sim 10$ pixels to facilitate removal of cosmic rays and CCD defects during the data reduction.
 The net effect of such dithering is to reduce the full exposure field of view to less than the full field of the CCD., The net effect of such dithering is to reduce the full exposure field of view to less than the full field of the CCD.
" All images used in this aualvsis have EWIIM sccing between L.7 to 2.6 pixels. or 2.9"" to L1"". somewhat larger than the typical seeiug at the Ningloug observing station."," All images used in this analysis have FWHM seeing between 1.7 to 2.6 pixels, or $''$ to $''$, somewhat larger than the typical seeing at the Xinglong observing station."
 Diving the observations the gain of the CCD system was adjusted several times (for various reasons Owing to having the CCD svstem work well). resulting iu gaius of 3.7 /aduiu 1995. 11 Jadu in 1996. aud 3.3 e faduin 1997.," During the observations the gain of the CCD system was adjusted several times (for various reasons owing to having the CCD system work well), resulting in gains of 3.7 $^{-}$ /adu in 1995, 4.1 $^{-}$ /adu in 1996, and 3.3 $^{-}$ /adu in 1997."
 Readout noise was coustaut (12 .) during the three years., Readout noise was constant (12 $^{-}$ ) during the three years.
 The mean value of overscau-subtracted bias was stable within auv one month period of time., The mean value of overscan-subtracted bias was stable within any one month period of time.
 LO bias frames (5 at the start of the night. 5 at the end) are taken daily for the BATC program.," 10 bias frames (5 at the start of the night, 5 at the end) are taken daily for the BATC program."
 As we can track the stability of the bias during the vear. we are able to average from 200 to 300 individual bias frames for cach welt of observation. which removes anv stable structure in the bias frame (amv variable level is removed via the overscan subtraction).," As we can track the stability of the bias during the year, we are able to average from 200 to 300 individual bias frames for each night of observation, which removes any stable structure in the bias frame (any variable level is removed via the overscan subtraction)."
 It is this average of niuiy. bias frames that is subtracted from the images of a given nieht., It is this average of many bias frames that is subtracted from the images of a given night.
 The CCD dark count has been constantly monitored throughout the BATC survey. iud las always been found to be stable aud rather free of gradients.," The CCD dark count has been constantly monitored throughout the BATC survey, and has always been found to be stable and rather free of gradients."
" As the average dark count/pixel is ο ο /hour. the dark count level in a 12.5 hour exposure is 128 /pixel. or of the sky level in this combined inaege and the spatial variation of dark is even smaller: Ίνοι, of negligible inportauce."," As the average dark count/pixel is 3 $^-$ /hour, the dark count level in a 42.8 hour exposure is 128 $^-$ /pixel, or of the sky level in this combined image and the spatial variation of dark is even smaller; i.e., of negligible importance."
 This coustaut dark count value was scaled to exposure time for cach image. then subtracted.," This constant dark count value was scaled to exposure time for each image, then subtracted."
 As detailed in our previous papers (cf., As detailed in our previous papers (cf.
 Fan et al., Fan et al.
 1996: Zhene et al., 1996; Zheng et al.
 1999). the BATC program las developed the means by which we can use dome fats to obtain accurate. high signal-to-noise (S/N). flat fields for flat-ficlding sky iuages.," 1999), the BATC program has developed the means by which we can use dome flats to obtain accurate, high signal-to-noise (S/N), flat fields for flat-fielding sky images."
 Briefly. the Ninglone Sclunidt telescope is equipped with a UV-trausparent plastic diffuser plate that can be firmly placed directly in front of the Schuudt corrector.," Briefly, the Xinglong Schmidt telescope is equipped with a UV-transparent plastic diffuser plate that can be firmly placed directly in front of the Schmidt corrector."
 The diffuser, The diffuser
"representation of the parent cloud but do show the cores in absorption as seen in the observations,",representation of the parent cloud but do show the cores in absorption as seen in the observations.
 With the possible exception of G31.03+00.26 Core B (which may have evidence of an 8m point source). the um images shows no emission in either the model or the observed data. as IRDCs do not emit significantly in the MIR due to their low temperatures.," With the possible exception of G31.03+00.26 Core B (which may have evidence of an $\mu$ m point source), the $\mu$ m images shows no emission in either the model or the observed data, as IRDCs do not emit significantly in the MIR due to their low temperatures."
 The modelled images show emission starting at gm for GO31.03+00.26 Core A and um for the remaining five cores (albeit only faintly for the cores in GO31.03+00.76) and absorption at shorter wavelengths., The modelled images show emission starting at $\mu$ m for G031.03+00.26 Core A and $\mu$ m for the remaining five cores (albeit only faintly for the cores in G031.03+00.76) and absorption at shorter wavelengths.
 The modelled emission fits well with the observations. at all wavelengths.," The modelled emission fits well with the observations, at all wavelengths."
 The masses of the cores were calculated in two ways: (1) using the radiative transfer models: and (11) using the observed flux densities at 4m. The latter was carried out using the equation (2).. whereBATky: is the source distance.," The masses of the cores were calculated in two ways: (i) using the radiative transfer models; and (ii) using the observed flux densities at $\mu$ m. The latter was carried out using the equation \citep{hildebrand83}, where is the source distance."
 We call the former the model mass and the latter the SED mass., We call the former the model mass and the latter the SED mass.
 Two different values of opacity at gam were. used: Kz0.03 cen gg! (?) and «20.05 ceni ge! (?2)..," Two different values of opacity at $\mu$ m were used: $\kappa$ $^{2}$ $^{-1}$ \citep{ossenkopf94} and $\kappa$ $^{2}$ $^{-1}$ \citep{preibisch93, henning95}."
 The latter is the value used to calculate the masses of low-mass prestellar cores (e.g. 2))., The latter is the value used to calculate the masses of low-mass prestellar cores (e.g. \citealt{kirk05}) ).
 The model masses were found using the ? opacities., The model masses were found using the \citet{ossenkopf94} opacities.
 The model masses match well with the SED masses calculated using κ-θιθδοσπ gg!.. but not with those calculated using «=0.03 cen gg!.," The model masses match well with the SED masses calculated using $\kappa$ $^{2}$ $^{-1}$, but not with those calculated using $\kappa$ $^{2}$ $^{-1}$."
 Both sets of SED masses and the model masses are shown in Table I.., Both sets of SED masses and the model masses are shown in Table \ref{coreprop}.
 The masses calculated here are in agreement with those found in previous studies of IRDCs (e.g. 222)).," The masses calculated here are in agreement with those found in previous studies of IRDCs (e.g. \citealt{williams01, simon06b, rathborne06}) )."
 In particular. in ?.. masses of dark cores were calculated and found to vary from 10 to 2100M...," In particular, in \citet{rathborne06}, masses of dark cores were calculated and found to vary from 10 to ${_\odot}$."
 All of our cores fit within this range., All of our cores fit within this range.
 8504m masses for both cores in. GO31.03+00.26 were calculated by ?.., $\mu$ m masses for both cores in G031.03+00.26 were calculated by \citet{parsons09}.
 Core A was found to have a mass of ΜΜ. and Core B of MM..., Core A was found to have a mass of $_{\odot}$ and Core B of $_{\odot}$.
 When using the same core boundaries as in? and «20.05 cen gg!. we find 500m masses of MM. and MM... similar to the 850gm masses.," When using the same core boundaries as in \citet{parsons09} and $\kappa$ $^{2}$ $^{-1}$, we find $\mu$ m masses of $_{\odot}$ and $_{\odot}$, similar to the $\mu$ m masses."
 This gives us additional confidence in our calculations., This gives us additional confidence in our calculations.
 The energy deposited in the gas due to cosmic ray tonisation heating is much lower than the energy absorbed by the dust due to external heating from the ISRF (?).., The energy deposited in the gas due to cosmic ray ionisation heating is much lower than the energy absorbed by the dust due to external heating from the ISRF \citep{evans01}.
 Thus. we can safely assume that the dust temperature is not affected by the energy exchange between gas and dust.," Thus, we can safely assume that the dust temperature is not affected by the energy exchange between gas and dust."
 The lack of MIR point sources also implies that heating via compression 1s minimal (2).., The lack of MIR point sources also implies that heating via compression is minimal \citep{battersby10}.
 These cores (again with the possible exception of GO031.03+00.26 Core B) therefore have no significant internal heating sources and are heated mainly by the external radiation field., These cores (again with the possible exception of G031.03+00.26 Core B) therefore have no significant internal heating sources and are heated mainly by the external radiation field.
 The temperature profiles of the modelled cores are shown in Fig 5.., The temperature profiles of the modelled cores are shown in Fig \ref{temprad}.
 They show a low temperature throughout the centre of the core. before it eventually begins rising. peaking at the outer edges.," They show a low temperature throughout the centre of the core, before it eventually begins rising, peaking at the outer edges."
 A higher temperature at the edges than at the centre is expected due to the lack of any internal heating source., A higher temperature at the edges than at the centre is expected due to the lack of any internal heating source.
 Some stepping can be seen in the graphs. this is due to the temperature table having a spacing of KK. This gradient is. consistent with previous studies of IRDCs (e.g. ?)) and is similar to what 1s seen in low-mass prestellar cores (e.g. 2)).," Some stepping can be seen in the graphs, this is due to the temperature table having a spacing of K. This gradient is consistent with previous studies of IRDCs (e.g. \citealt{peretto10}) ) and is similar to what is seen in low-mass prestellar cores (e.g. \citealt{nutter08}) )."
 However. IRDCs have higher outer temperatures than low-mass prestellar cores. which can be explained by a higher ISRF around these IRDCs relative to nearby regions of star formation.," However, IRDCs have higher outer temperatures than low-mass prestellar cores, which can be explained by a higher ISRF around these IRDCs relative to nearby regions of star formation."
 As all six cores have a flattened geometry. the temperature-radius profile varies with the direction inside the core.," As all six cores have a flattened geometry, the temperature-radius profile varies with the direction inside the core."
 The two extreme cases. 9=O° and &=90°. are shown in the temperature-radius profiles.," The two extreme cases, $\theta=0\degr$ and $\theta=90\degr$, are shown in the temperature-radius profiles."
 €=90° is along the midplane of the structure. while €.=0° corresponds to the short axis of a flattened core. perpendicular to the midplane.," $\theta=90\degr$ is along the midplane of the structure, while $\theta=0\degr$ corresponds to the short axis of a flattened core, perpendicular to the midplane."
 The average temperature from the model of each core agrees with the result from the single-temperature greybody fit., The average temperature from the model of each core agrees with the result from the single-temperature greybody fit.
 However. the temperature-radius profile shows this fit to be an over-simplification. with temperatures within each core varying by - KK and being dependent upon angular direction in the core.," However, the temperature-radius profile shows this fit to be an over-simplification, with temperatures within each core varying by $\sim$ K and being dependent upon angular direction in the core."
 Five of the six cores required a higher ISRF than the solar neighbourhood value in order to correctly match the observed flux densities at shorter wavelengths (ie. 160 and jm)., Five of the six cores required a higher ISRF than the solar neighbourhood value in order to correctly match the observed flux densities at shorter wavelengths (i.e. 160 and $\mu$ m).
 As an example we show the cores’ ISRF at far-ultraviolet (FUV) in Table 1.., As an example we show the cores' ISRF at far-ultraviolet (FUV) in Table \ref{coreprop}.
 This is merely to indicate how the ISRF varies from core to core., This is merely to indicate how the ISRF varies from core to core.
 For comparison. the solar neighbourhood [SRF at FUV is 7.6x107 eere s! em (2). or 4.6 times the Habing flux (1.6:107 eere s! em. 2).," For comparison, the solar neighbourhood ISRF at FUV is $\times$ $^{-3}$ erg $^{-1}$ $^{-2}$ \citep{black94} or 4.6 times the Habing flux $\times10^{-3}$ erg $^{-1}$ $^{-2}$, \citealt{habing68}) )."
 The IRDCs typically have an ISRF around four times greater. consistent with them being in the stronger interstellar radiation field of the inner Galactic Plane.," The IRDCs typically have an ISRF around four times greater, consistent with them being in the stronger interstellar radiation field of the inner Galactic Plane."
 Differences between the ISRF required for cores within the same field can be explained by one core being more deeply buried. within the parent cloud and thus having less external radiation heating it., Differences between the ISRF required for cores within the same field can be explained by one core being more deeply buried within the parent cloud and thus having less external radiation heating it.
 The difference between the ISRFs of the cores are listed in Table | as /nο. where { corresponds to the core that is modelled using a (+).higher ISRF and 7» to the core modelled with the lower ISRF (within the same field).," The difference between the ISRFs of the cores are listed in Table \ref{coreprop} as $ln\left(\frac{I_1}{I_2}\right)$, where $I_1$ corresponds to the core that is modelled using a higher ISRF and $I_2$ to the core modelled with the lower ISRF (within the same field)."
 This ratio can be thought às à measure of how more deeply a core is embedded when compared with the other core in the same field., This ratio can be thought as a measure of how more deeply a core is embedded when compared with the other core in the same field.
 Core B in GO31.03+00.26 appears to reside in an unusually intense radiation field. 17 times that of the solar neighbourhood and far higher thàn any of the other cores. possibly pointing to an abnormally high ISRF in this field.," Core B in G031.03+00.26 appears to reside in an unusually intense radiation field, 17 times that of the solar neighbourhood and far higher than any of the other cores, possibly pointing to an abnormally high ISRF in this field."
 Core A in the same cloud has an ISRF value comparable to the other cores. which implies that it is either more deeply buried within the IRDC than Core B or there is a highly asymmetrical external," Core A in the same cloud has an ISRF value comparable to the other cores, which implies that it is either more deeply buried within the IRDC than Core B or there is a highly asymmetrical external"
signal lies the binning process as well as the summation over angular frequency bins and redshift bins.,signal lies the binning process as well as the summation over angular frequency bins and redshift bins.
" Since the signs of the biases contributed by different angular frequencies and redshifts can be different, there can be bias cancellation during these processes."," Since the signs of the biases contributed by different angular frequencies and redshifts can be different, there can be bias cancellation during these processes."
" In another word, |b;/b;| can depend on binning choices."," In another word, $|\rm{b}_f/\rm{b}_i|$ can depend on binning choices."
" We then vary two cosmological parameters (Ωμ, and og) and investigate how the original and the nulled parameter constraints change with respect to the number of redshift bins available.", We then vary two cosmological parameters $\Omega_{\rm m}$ and $\sigma_{\rm 8}$ ) and investigate how the original and the nulled parameter constraints change with respect to the number of redshift bins available.
 For this case we use all triangle shapes to enable a comparison with results for two-point statistics., For this case we use all triangle shapes to enable a comparison with results for two-point statistics.
" Our result 5) shows that a further increase of the number of redshift bins beyond 10 is not very rewarding in terms of information preservation as characterized by the FoM, in either 2p, 3p, or 2p+3p cases."," Our result $\,$ ) shows that a further increase of the number of redshift bins beyond 10 is not very rewarding in terms of information preservation as characterized by the FoM, in either 2p, 3p, or 2p+3p cases."
" This suggests, when the possibility of more redshift bins exists, the choice of redshift bin number should be based mainly on the requirement of bias reduction level in case of negligible photometric errors."," This suggests, when the possibility of more redshift bins exists, the choice of redshift bin number should be based mainly on the requirement of bias reduction level in case of negligible photometric errors."
" When there are non-negligible photometric errors, however, the information loss will probably be more severe, as found by for the two-point case."," When there are non-negligible photometric errors, however, the information loss will probably be more severe, as found by for the two-point case."
 The necessity of carrying out data compression in cosmology has long been recognized and has been ever increasing due to the increasing size of the data sets., The necessity of carrying out data compression in cosmology has long been recognized and has been ever increasing due to the increasing size of the data sets.
 In cosmic shear studies the survey area of next generation multicolor imaging surveys will be an order of magnitude larger than the current ones., In cosmic shear studies the survey area of next generation multicolor imaging surveys will be an order of magnitude larger than the current ones.
" The study of three-point statistics also implies a huge increase in the amount of data directly entering the Fisher-matrix/ likelihood analysis, compared to the two-point case."," The study of three-point statistics also implies a huge increase in the amount of data directly entering the Fisher-matrix/ likelihood analysis, compared to the two-point case."
Recent observational evidence of ubiquitous periodically varying features in the solar corona (e.g..Tomezyketal.2007;2009) has raised the debate on whether these observations are caused by magnetohydrodynamiec (MHD) waves or by quasi-periodic flows (see.e.g..DePontieu&McIntosh2010).,"Recent observational evidence of ubiquitous periodically varying features in the solar corona \citep[e.g.,][]{tomczyk2007,jess2009,tomczyk2009,wang2009} has raised the debate on whether these observations are caused by magnetohydrodynamic (MHD) waves or by quasi-periodic flows \citep[see, e.g.,][]{depontieu2010}."
. There seem to be strong theoretical arguments supporting the wave interpretation (e.g..Erdélyi&Fedun2007:VanDoorsselaereetal.2008:Terradas2010:Verth 2010).," There seem to be strong theoretical arguments supporting the wave interpretation \citep[e.g.,][]{erdelyi2007,tom2008,TGV,VTG2010}."
. However. waves and flows are not mutually exclusive and. in fact. both phenomena have been simultaneously observed in the fine structure. of solar prominences (e.g..Okamotoetal.," However, waves and flows are not mutually exclusive and, in fact, both phenomena have been simultaneously observed in the fine structure of solar prominences \citep[e.g.,][]{okamoto}."
2007).. This offers us the opportunity to study the interaction between waves and flows in the solar atmosphere., This offers us the opportunity to study the interaction between waves and flows in the solar atmosphere.
 The fine structure of solar prominences is clearly visible in the high-resolution Ha and Ca II H-line images from the Solar Optical Telescope (SOT) aboard the Hinode satellite (e.g..Okamotoetal.2007;Berger2008;Chae2008:Ningetal.2009;Schmieder2010:Chae 2010).," The fine structure of solar prominences is clearly visible in the high-resolution $\alpha$ and Ca II H-line images from the Solar Optical Telescope (SOT) aboard the Hinode satellite \citep[e.g.,][]{okamoto,berger,chae,ning,brigi,chae2010}."
. When observed above the limb. vertical structures are commonly seen in quiescent prominences (e.g..Bergeretal.2008:Chae 2010).. while horizontal threadlike structures are usually observed in active region prominences (e.g..Okamotoetal.," When observed above the limb, vertical structures are commonly seen in quiescent prominences \citep[e.g.,][]{berger,chae,chae2010}, while horizontal threadlike structures are usually observed in active region prominences \citep[e.g.,][]{okamoto}."
 2007).. Although it is apparently difficult to reconcile both pictures. some authors (e.g..Schmiederetal.2010) have suggested that vertical threads might actually be a pile up of horizontal threads which appear às vertical structures when projected on the plane of the sky.," Although it is apparently difficult to reconcile both pictures, some authors \citep[e.g.,][]{brigi} have suggested that vertical threads might actually be a pile up of horizontal threads which appear as vertical structures when projected on the plane of the sky."
 This idea is consistent with H« observations of filaments on the solar disk from the Swedish Solar Telescope (e.g..Linetal.2008.2009).. in which the filament fine structure is seen as thin and long dark ribbons.," This idea is consistent with $\alpha$ observations of filaments on the solar disk from the Swedish Solar Telescope \citep[e.g.,][]{lin08,lin09}, in which the filament fine structure is seen as thin and long dark ribbons."
 On the other hand. other authors (e.g..Chae2010) have argued that vertical threads are real and are an indication of the existence of vertical magnetic fields in quiescent prominences.," On the other hand, other authors \citep[e.g.,][]{chae2010} have argued that vertical threads are real and are an indication of the existence of vertical magnetic fields in quiescent prominences."
 Thus. it remains unclear whether all prominences have the same magnetic structure or. on the contrary. the magnetic field in quiescent prominences 1s predominantly vertical and active region prominences have horizontal fields.," Thus, it remains unclear whether all prominences have the same magnetic structure or, on the contrary, the magnetic field in quiescent prominences is predominantly vertical and active region prominences have horizontal fields."
 A recent review on the properties of prominence threads can be found in Lin(2010).., A recent review on the properties of prominence threads can be found in \citet{linrev}.
 There are many evidences of transverse oscillations of the fine structures of both active region and quiescent prominences. which have been interpreted in terms of kink MHD waves (seeBallester 2010)..," There are many evidences of transverse oscillations of the fine structures of both active region and quiescent prominences, which have been interpreted in terms of kink MHD waves \citep[see the recent reviews by][]{ballester, oliver, arreguiballester}."
 The reported periods are usually in a narrow band between 2 and 10 minutes. while the oscillations are typically damped after a few periods.," The reported periods are usually in a narrow band between 2 and 10 minutes, while the oscillations are typically damped after a few periods."
 In addition. flows and mass motions in prominences have been also reported (e.g..al.2003:Ahnet 2010).," In addition, flows and mass motions in prominences have been also reported \citep[e.g.,][]{zirker98, wang1999, kucera2003, lin2003,ahn2010}."
 The typical flow velocities are less than 30 km s! in quiescent prominences. although larger values up to 40—50 km s~! have been observed in active region prominences.," The typical flow velocities are less than 30 km $^{-1}$ in quiescent prominences, although larger values up to 40–50 km $^{-1}$ have been observed in active region prominences."
 The work of Okamotoetal.(2007) is an example of simultaneous transverse oscillations and mass flows in prominence fine structures observed with Hinode/SOT., The work of \citet{okamoto} is an example of simultaneous transverse oscillations and mass flows in prominence fine structures observed with Hinode/SOT.
 In the present paper we focus on the theoretical analysis of the event reported by Okamotoetal.(2007)., In the present paper we focus on the theoretical analysis of the event reported by \citet{okamoto}.
. Similar observations of simultaneous flows and oscillations have been reported by Ofman&Wang(2008) in coronal loops and by Cao(2010) in filament footpoints., Similar observations of simultaneous flows and oscillations have been reported by \citet{ofmanwang} in coronal loops and by \citet{cao2010} in filament footpoints.
 Also. the recent work by Antolit&Verwichte(2011) on observations of transverse oscillations of loops with coronal rain is relevant for our present theoretical investigation.," Also, the recent work by \citet{antolin} on observations of transverse oscillations of loops with coronal rain is relevant for our present theoretical investigation."
 Okamotoetal.(2007) observed an active regior prominence formed by a myriad of horizontal magnetic flux tubes which are partially outlined bythreads of cool and dense prominence plasma., \citet{okamoto} observed an active region prominence formed by a myriad of horizontal magnetic flux tubes which are partially outlined bythreads of cool and dense prominence plasma.
 The magnetic tubes are probably rooted in the solar photosphere., The magnetic tubes are probably rooted in the solar photosphere.
 Although only the part of the tubes filled with prominence material can be seen in the Ca II H- images. the length of the whole magnetic tube must be much longer than the length of the prominence threads. which is roughly between 3.000 km and 16.000 km.," Although only the part of the tubes filled with prominence material can be seen in the Ca II H-line images, the length of the whole magnetic tube must be much longer than the length of the prominence threads, which is roughly between 3,000 km and 16,000 km."
 Okamotoetal.(2007) detected that some threads were flowing along the magnetic tubes and simultaneously oscillating in the vertical direction., \citet{okamoto} detected that some threads were flowing along the magnetic tubes and simultaneously oscillating in the vertical direction.
 The mean period of the oscillations was 3 min and the apparent flow velocity on the plane of sky was around 40 km s7!., The mean period of the oscillations was 3 min and the apparent flow velocity on the plane of sky was around 40 km $^{-1}$ .
 The oscillations were in phase along the whole length of the threads. and the wavelength was estimated to be at least 250.000 km.," The oscillations were in phase along the whole length of the threads, and the wavelength was estimated to be at least 250,000 km."
»v an external mass distribution. may be mocdoelled by adeding a shear term to the lens equation).,by an external mass distribution may be modelled by adding a shear term to the lens equation).
 However. for an investigation of the lensing due to large scale structure. where the mass distribution is assumed to be continuous and 1o0miogeneous in all directions about à ray. we may introduce an artificial shear on rays near the shooting boundary.," However, for an investigation of the lensing due to large scale structure, where the mass distribution is assumed to be continuous and homogeneous in all directions about a ray, we may introduce an artificial shear on rays near the shooting boundary."
 For the RBAL we choose to calculate each of he deflection angles explicitlv. with an increase in the computational time over an equivalent. hierarchical method.," For the RBM, we choose to calculate each of the deflection angles explicitly, with an increase in the computational time over an equivalent hierarchical method."
 Dv making the direct. caleulation of the dellection. we are ree to choose the geometry of the region within which we include lenses.," By making the direct calculation of the deflection, we are free to choose the geometry of the region within which we include lenses."
 The most natural choice for a distribution which is homogeneous and. isotropic beyond. some length scale fy is to include lenses within a circular region around each rav out to the μασας Rien.=£g., The most natural choice for a distribution which is homogeneous and isotropic beyond some length scale $R_{\rm H}$ is to include lenses within a circular region around each ray out to the radius $R_{\rm lens} = R_{\rm H}$.
 For the isolated lens geometries we examine here. we need only set Rieu. to encompass the RSAL image plane.," For the isolated lens geometries we examine here, we need only set $R_{\rm lens}$ to encompass the RSM image plane."
 IX. discussion of the appropriate choice for Ay in cosmological lensing scenarios is reserved for Fluke et al. (, A discussion of the appropriate choice for $R_{\rm H}$ in cosmological lensing scenarios is reserved for Fluke et al. (
in preparation).,in preparation).
 As a final test. here. we now consider a random. distribution of Mia.=5 lenses (with the same lens positions or both the RSAL and RBA.," As a final test here, we now consider a random distribution of $N_{\rm lens} = 5$ lenses (with the same lens positions for both the RSM and RBM)."
 The. resulting MILL. is shown in ancl various parameters in2., The resulting MPH is shown in and various parameters in.
". For xh. methods we use a total of 10"". bundles/sources. but discarding bundles within rea=LOL Einstein radii for the UDM reduces this to 9.310 bundles for the histogram."," For both methods we use a total of $10^6$ bundles/sources, but discarding bundles within $x_{\rm cut} = 1.01$ Einstein radii for the RBM reduces this to $9.3\times10^5$ bundles for the histogram."
 The sample mean and variance are Gr)=1.09. ση=2.09 or the RBAL and (560=0.98. Lh=0.23 for the ΛΙ.," The sample mean and variance are $\langle\mu\rangle = 1.09$, $\sigma^2_{\mu} = 2.09$ for the RBM and $\langle\mu\rangle = 0.98$, $\sigma^2_{\mu} = 0.23$ for the RSM."
 Using an ensemble of lenses introduces a new length scale to the problem. so that sub-structure at intermeciate magnifications (32jrz 30) is seen in the MPIL using he RSAL (but not with the ΔΕΟ.," Using an ensemble of lenses introduces a new length scale to the problem, so that sub-structure at intermediate magnifications $3 \simeq \mu \simeq 30$ ) is seen in the MPH using the RSM (but not with the RBM)."
 The ‘hump’ in the AIPIL occurs for a planar distribution of lenses. and. has xen studied: both numerically (Rauchetal.1992). and analytically (Ixofmanetal.1997)... and is believed to be due o the caustic patterns of pairs of point lenses.," The `bump' in the MPH occurs for a planar distribution of lenses, and has been studied both numerically \cite{RAUCH92} and analytically \cite{KOFMAN97}, and is believed to be due to the caustic patterns of pairs of point lenses."
 Lf this is a caustic-incuced feature. ic.," If this is a caustic-induced feature, ie."
 a region of high magnilication. hen the RM will not. provide the ‘correct’ magnification.," a region of high magnification, then the RBM will not provide the `correct' magnification."
 We [eel the feature may be in part due to the low resolution with whieh the complex caustic structure was mapped with he RSAL on a regular grid (10001000 pixels). however a discussion of this elfect is beyond the scope of this paper.," We feel the feature may be in part due to the low resolution with which the complex caustic structure was mapped with the RSM on a regular grid $1000 \times 1000$ pixels), however a discussion of this effect is beyond the scope of this paper."
 The Ray Bundle method. provides a computationally fast. accurate and Blexible. alternative to the Ray Shooting method for studies of the weak gravitational lensing limit.," The Ray Bundle method provides a computationally fast, accurate and flexible alternative to the Ray Shooting method for studies of the weak gravitational lensing limit."
 A wide variety of lens models are easily incorporated ο have considered here only the case of the Schwarzschilel lens. but changing to a cilferent mocel involves mocdifving only a(&) in the GLE.," A wide variety of lens models are easily incorporated – we have considered here only the case of the Schwarzschild lens, but changing to a different model involves modifying only $\hat{\bvec{\alpha}}(\bvec{\xi})$ in the GLE."
 One alternative to ray based methods requires solving for the scalar lensing potential. but it may not always be possible to find an analytic solution for all lens models.," One alternative to ray based methods requires solving for the scalar lensing potential, but it may not always be possible to find an analytic solution for all lens models."
 The RAM also allows us to avoid artificial shear introduced bv erid based methods for light ravs which pass near the shooting boundary. when a small portion of an otherwise homogeneous and isotropic distribution is used.," The RBM also allows us to avoid artificial shear introduced by grid based methods for light rays which pass near the shooting boundary, when a small portion of an otherwise homogeneous and isotropic distribution is used."
 The point source limit may be approached as values of I may be selected without the restriction of introducing a, The point source limit may be approached as values of $R_{\rm i}$ may be selected without the restriction of introducing a
surface gravity measurements ο[ος by 0.5 dex.,surface gravity measurements differ by 0.5 dex.
 This is most likely clue to a systematic ellect in. the mociel spectra themselves. with the Hla. line. profile analysis overestimating the surface eravity relative to an analysis involving the complete series., This is most likely due to a systematic effect in the model spectra themselves with the $\alpha$ line profile analysis overestimating the surface gravity relative to an analysis involving the complete series.
 The shape and strength. of the upper Balmer lines are very sensitive to surface eravity and oller a more reliable surface. gravity diagnostics., The shape and strength of the upper Balmer lines are very sensitive to surface gravity and offer a more reliable surface gravity diagnostics.
 In suninaryv we estimate the parameters. representative of the sdB GALEN J03211 4727. at phase 0.5., In summary we estimate the parameters representative of the sdB GALEX $+$ 4727 at phase 0.5.
 The effective temperature anc helium abundance. were estimated by taking the weighted average of the results from the LHoa- Η11 and Ho (0.35 - 0.65) fits and for the surface. gravity we adopted the result from the ΤΑΞΙ] analysis: and Applying a correction of. 0.5+0.2 to the surface gravity measurement of GALEN | 3844 that is based on Ia alone we conservatively estimate the sclB parameters: and Our analysis assumes a hycdrogen/helium composition anc the inclusion. of heavy elements. in. the πιο atmospheres is likely to allect our results (seeHeber.etal.2000:Edelmannet 2003)..," The effective temperature and helium abundance were estimated by taking the weighted average of the results from the $\alpha$ -H11 and $\alpha$ (0.35 - 0.65) fits and for the surface gravity we adopted the result from the $\alpha$ -H11 analysis: and Applying a correction of $-0.5\pm0.2$ to the surface gravity measurement of GALEX $+$ 3844 that is based on $\alpha$ alone we conservatively estimate the sdB parameters: and Our analysis assumes a hydrogen/helium composition and the inclusion of heavy elements in the model atmospheres is likely to affect our results \citep[see][]{heb2000,ede2003}."
 A self consistent analysis including abundance measurements (c.g..OLoole&Leber2006:Ahmadctal.2007). awaits high signal-to-noise and resolution ultraviolet and optical spectroscopy.," A self consistent analysis including abundance measurements \citep[e.g.,][]{oto2006,ahm2007} awaits high signal-to-noise and resolution ultraviolet and optical spectroscopy."
 Using the mass function (lable 4)) and assuming a mass of 0.5AL. for the hot. subclwarl GALEN | 472. we calculate a minimum secondary mass of 0.13.Al. that corresponds to a spectral type of ALS Osirkpatrick&Ale-Carthy 1904).," Using the mass function (Table \ref{tbl_bin}) ) and assuming a mass of $0.5\ M_\odot$ for the hot subdwarf GALEX $+$ 472, we calculate a minimum secondary mass of $0.13\ M_\odot$ that corresponds to a spectral type of M5 \citep{kir1994}."
.. Assuming an absolute J magnitude of 3.9 or the hot subclwark a ALS star with A;=8.8 would he outshone by the hot subchwarl.," Assuming an absolute $J$ magnitude of 3.9 for the hot subdwarf, a M5 star with $M_J = 8.8$ would be outshone by the hot subdwarf."
 The NSVS lighteurve shows C:ALEN | 4727 to »* variable ancl the search for a period in the photonietric data resulted in a best-period corresponding to the orbital »eriod., The NSVS lightcurve shows GALEX $+$ 4727 to be variable and the search for a period in the photometric data resulted in a best-period corresponding to the orbital period.
 The observed. variations are most probably caused w irradiation of the atmosphere of the cool companion by he hot subcdwarf., The observed variations are most probably caused by irradiation of the atmosphere of the cool companion by the hot subdwarf.
 Using the observed. semi-amplitude of 061 mag we may constrain the binary. parameters further., Using the observed semi-amplitude of 0.061 mag we may constrain the binary parameters further.
 We estimated the svstem inclination and secondary mass using the rellection. model of Alaxtecdetal.(2002). andassuming two cillerent. masses for the κα star., We estimated the system inclination and secondary mass using the reflection model of \citet{max2002} andassuming two different masses for the sdB star.
 For a sdB mass of 0.4 AZ... the inclination is predicted to be between," For a sdB mass of $0.4\ M_\odot$ , the inclination is predicted to be between"
reported in the literature that include radiation feedback in the context of the typical region of Galactic star cluster formation. ax opposed to focusing on single low-mass (?) or high-1uass (7277). cores. or on low-1ass or low-density regions like Taurus (277)..,"reported in the literature that include radiation feedback in the context of the typical region of Galactic star cluster formation, as opposed to focusing on single low-mass \citep{commercon10a} or high-mass \citep{krumholz07a, krumholz10a, myers11a} cores, or on low-mass or low-density regions like Taurus \citep{bate09a, offner10a, peters10b}."
 Our simulations returu a surprising result., Our simulations return a surprising result.
 At carly times in the simulations. accreting stars produce bubbles of warin. raciatively-heated gas around themselves. aud within these bubbles fragmentation is suppressed by the increased Bounor-Ehert mass.," At early times in the simulations, accreting stars produce bubbles of warm, radiatively-heated gas around themselves, and within these bubbles fragmentation is suppressed by the increased Bonnor-Ebert mass."
 However. we find that. once ~LO20% of the eas in the protocluster has been converted to stars. these bubbles of warm gas begin to overlap and imerec.," However, we find that, once $\sim 10-20\%$ of the gas in the protocluster has been converted to stars, these bubbles of warm gas begin to overlap and merge."
 Rather than resemibliug a few wari islands surouuded by a sea of cold gas. we instead have a cloud where all the gas is warmed by the collective 1uninositv of all the accreting stars.," Rather than resembling a few warm islands surrounded by a sea of cold gas, we instead have a cloud where all the gas is warmed by the collective luminosity of all the accreting stars."
 Ouce the simulation reaches this state. vacation ecdhack raises the temperature and the Bounor-Ebert nass throughout the remaining gas enough to esseutiallv ια the formation of anv further stars.," Once the simulation reaches this state, radiation feedback raises the temperature and the Bonnor-Ebert mass throughout the remaining gas enough to essentially halt the formation of any further stars."
 Mass continues o be converted from gas to stars. but this is almost entirely through accretion outo existing stars rather than ornmation of new ones.," Mass continues to be converted from gas to stars, but this is almost entirely through accretion onto existing stars rather than formation of new ones."
 As a result. when radiation is included. the stellar mass distribution iu a elobally-collapsing star cluster such as the one we simulate is not jearlv constaut or very slightly iucreasing with tine. as as been reported iu earlier. non-radiative simulations. and as we fud here in a control run that does nof include radiation.," As a result, when radiation is included, the stellar mass distribution in a globally-collapsing star cluster such as the one we simulate is not nearly constant or very slightly increasing with time, as has been reported in earlier, non-radiative simulations, and as we find here in a control run that does not include radiation."
 Instead. the stellar mass distribution shifts strongly to systematically higher masses as star orlation proceeds. eventually becoming too top-heavy compared to the observed IME.," Instead, the stellar mass distribution shifts strongly to systematically higher masses as star formation proceeds, eventually becoming too top-heavy compared to the observed IMF."
 While the absolute mass scale remains uncertain in our simulations due to our inability to resolve tight binaries. the result that the IME Is non-constaut and increasing with time is robust agaiust changes iu resolution.," While the absolute mass scale remains uncertain in our simulations due to our inability to resolve tight binaries, the result that the IMF is non-constant and increasing with time is robust against changes in resolution."
 This unuplies that. unless there is also some niechanisni to ensure that star formation in every protocluster stops when the IME peak is in the sale place. it is nof possible to produce the invariaut ΤΠΕ peak that we observe via the elobal collapse scenario we have simulated.," This implies that, unless there is also some mechanism to ensure that star formation in every protocluster stops when the IMF peak is in the same place, it is not possible to produce the invariant IMF peak that we observe via the global collapse scenario we have simulated."
 We arene that the underline reason that this problem occurs is that. iu the absence of either external turbuleut dviving or any sort of internal mechanical feedback to slow star formation down. stars in our sinulation form too quickly.," We argue that the underlying reason that this problem occurs is that, in the absence of either external turbulent driving or any sort of internal mechanical feedback to slow star formation down, stars in our simulation form too quickly."
 Since accretion huninositv produced as gas falls outo stars is what ultimately drives the temperature Increase in our snmulatious that shuts off fragmentation and leads to a top-heavy IME. the problem is likely to be alleviated iun simulations that include enough plivsies to obtain a low star formation rate similar to that observed in real star clusters.," Since accretion luminosity produced as gas falls onto stars is what ultimately drives the temperature increase in our simulations that shuts off fragmentation and leads to a top-heavy IMF, the problem is likely to be alleviated in simulations that include enough physics to obtain a low star formation rate similar to that observed in real star clusters."
 We are in the process of conducting such simulations now. aud will report on the results i future publications.," We are in the process of conducting such simulations now, and will report on the results in future publications."
 We thank R. Banerjee. €. Federrath. BR. Nlessen. AL Mac Low. T. Peters for helpful discussions.," We thank R. Banerjee, C. Federrath, R. Klessen, M. Mac Low, T. Peters for helpful discussions."
 This work was supported bv an Alfred SSloau Fellowship (MBREK): the NSF through erauts CAREER-0955300 CXBIS) and AST-0807739. (MBRIS). and AST-0908553 (CEM and RIK): NASA through ΑΤΕΡ eraut NNNOQARSIC. (RIK. CEM. and AIRS) and a Spitzer Space Telescope Theoretical Research Program erat (CEM and AIRIN): aud the US Departineut of Energy at LLNL under coutrast DE-AC52-07NA (RIS).," This work was supported by an Alfred Sloan Fellowship (MRK); the NSF through grants CAREER-0955300 (MRK) and AST-0807739 (MRK), and AST-0908553 (CFM and RIK); NASA through ATFP grant NNX09AK31G (RIK, CFM, and MRK) and a Spitzer Space Telescope Theoretical Research Program grant (CFM and MRK); and the US Department of Energy at LLNL under contrast DE-AC52-07NA (RIK)."
 Support for computer simulations was provided by an LRAC eraut from the NSF through Teragrid resources aud NASA through erauts from the ATFP and Spitzer Theory Progra., Support for computer simulations was provided by an LRAC grant from the NSF through Teragrid resources and NASA through grants from the ATFP and Spitzer Theory Program.
within the timespan of the simulation described in Fig.6..,within the timespan of the simulation described in \ref{figure_e3}.
 This is why the disk appears so depleted at the end., This is why the disk appears so depleted at the end.
 Conversely. for low-mass planets. the close encounter probability ts so low that many planetesimals keep following the secular dynamies until the end of the run.," Conversely, for low-mass planets, the close encounter probability is so low that many planetesimals keep following the secular dynamics until the end of the run."
 Therefore. they have enough time to generate a strong asymmetric clump.," Therefore, they have enough time to generate a strong asymmetric clump."
 We stress here that this clump is not due to any mean- resonance., We stress here that this clump is not due to any mean-motion resonance.
 There is thus no need for planet migration in this case and this may appear as an alternative scenario to mean-motion resonance for generating transient clumps., There is thus no need for planet migration in this case and this may appear as an alternative scenario to mean-motion resonance for generating transient clumps.
 Nevertheless. even with an Earth mass planet. these clumps do not last as long as the resonant clumps.," Nevertheless, even with an Earth mass planet, these clumps do not last as long as the resonant clumps."
 Planetesimals are not protected against close encounters with the planet. which finally deplete most of the disk after 35 Myr. in our simulations.," Planetesimals are not protected against close encounters with the planet, which finally deplete most of the disk after $35$ Myr, in our simulations."
 As explained in the previous sections. we have run many simulations in order to sample correctly the parameter space of planetary eccentricity versus planetary mass.," As explained in the previous sections, we have run many simulations in order to sample correctly the parameter space of planetary eccentricity versus planetary mass."
 We can thus address the question of the visibility of asymmetric structures in a disk., We can thus address the question of the visibility of asymmetric structures in a disk.
 As already mentioned. the shape of a disk i5 dominated either by resonant or by non-resonant planetesimals.," As already mentioned, the shape of a disk is dominated either by resonant or by non-resonant planetesimals."
 The region where MMRs dominate the disk shape correspond to planets on low-eecentricity orbits., The region where MMRs dominate the disk shape correspond to planets on low-eccentricity orbits.
 In. this region. two situations can occur: planets can generate clear resonant patterns with several visible clumps in the disk (generally while on circular or very low eccentric orbits) or produce smooth patterns. with only a hole at the planet location as the visible structure (generally while on an orbitwith an eccentricity between 0.05 and 0.1).," In this region, two situations can occur: planets can generate clear resonant patterns with several visible clumps in the disk (generally while on circular or very low eccentric orbits) or produce smooth patterns, with only a hole at the planet location as the visible structure (generally while on an orbitwith an eccentricity between $0.05$ and $0.1$ )."
 Concerning the non-resonant planetesimals. Earth mass planets on eccentric orbits can generate observable structures by secular perturbations.," Concerning the non-resonant planetesimals, Earth mass planets on eccentric orbits can generate observable structures by secular perturbations."
 Outside these regions. the disks do hot show any observable structures. when they are not totally depleted.," Outside these regions, the disks do not show any observable structures, when they are not totally depleted."
 The results for all these simulations are summarized in Table 2 and Fig. 10.., The results for all these simulations are summarized in Table \ref{normal} and Fig. \ref{resume}.
 Three main regions can be identified., Three main regions can be identified.
 In zones I and Il. observable structures in the disk are generated by MMRs while in zone IIL. transient structures are generated by non-resonant mechanisms.," In zones I and II, observable structures in the disk are generated by MMRs while in zone III, transient structures are generated by non-resonant mechanisms."
 In the remaining region. the disk does not show any structure.," In the remaining region, the disk does not show any structure."
 In zone Il. MMRs create clumpy disks while in zone II they generate a smooth disk with a hole at the planet location.," In zone I, MMRs create clumpy disks while in zone II they generate a smooth disk with a hole at the planet location."
 This figure also shows the fraction of planetesimals still bound to the system after 40 Myr (background color)., This figure also shows the fraction of planetesimals still bound to the system after $40$ Myr (background color).
 However. this quantity is sensitive to several parameters (stellar mass. duration of the simulation. initial distribution of the planetesimals ...) while the limits of the three zones are quite independent.," However, this quantity is sensitive to several parameters (stellar mass, duration of the simulation, initial distribution of the planetesimals ...) while the limits of the three zones are quite independent."
 However. we have assumed for these simulations a constant migration rate of 0.5 AU Myr! and a disk of planetesimals initially on circular orbits.," However, we have assumed for these simulations a constant migration rate of $0.5$ AU $^{-1}$ and a disk of planetesimals initially on circular orbits."
 With different assumptions. the outcomes of the simulations could be changed: we have thus investigated these two parameters in order to discuss the robustness of our conclusions., With different assumptions the outcomes of the simulations could be changed: we have thus investigated these two parameters in order to discuss the robustness of our conclusions.
 For a given planetary mass and eccentricity. we expect the structures to change. as the trapping probability depends on the migration rate and on the planetesimal eccentricity.," For a given planetary mass and eccentricity, we expect the structures to change, as the trapping probability depends on the migration rate and on the planetesimal eccentricity."
 But. again. our main focus ts to determinate if the resonant structures are visible or not.," But, again, our main focus is to determinate if the resonant structures are visible or not."
 For instance. in the low-eccentricity orbit case. Neptune mass and Jupiter mass planets do not produce the same structures but they have the same sensitivity to the planet eccentricity.," For instance, in the low-eccentricity orbit case, Neptune mass and Jupiter mass planets do not produce the same structures but they have the same sensitivity to the planet eccentricity."
 Here. we investigate if the migration rate or the initial planetesimal eccentricity change these conclusions.," Here, we investigate if the migration rate or the initial planetesimal eccentricity change these conclusions."
 The results of these additional simulations are summarized in Tables 3. and 4.. in the same manner as in Table 2. for the nominal case.," The results of these additional simulations are summarized in Tables \ref{fastslow} and \ref{complete_hot}, in the same manner as in Table \ref{normal} for the nominal case."
 Table 3 corresponds to simulations with different migration rates but unexcited initial planetesimal disks., Table \ref{fastslow} corresponds to simulations with different migration rates but unexcited initial planetesimal disks.
 Table 4 describes simulations with initially excited disks., Table \ref{complete_hot} describes simulations with initially excited disks.
" In our work. the migration rate parameter was chosen to match the best fit obtained by ? for the Vega debris disk to allow direct comparison,"," In our work, the migration rate parameter was chosen to match the best fit obtained by \citet{2003ApJ...598.1321W} for the Vega debris disk to allow direct comparison."
 In his paper. Wyatt discussed the impact of the migration rate in the restrictive case of a planet on a circular orbit. and has shown that the trapping probability in an MMR increases with decreasing migration rates (or when the star is less massive).," In his paper, Wyatt discussed the impact of the migration rate in the restrictive case of a planet on a circular orbit, and has shown that the trapping probability in an MMR increases with decreasing migration rates (or when the star is less massive)."
 We have thus performed simulations with a lower (0.05 AU Myr) and a higher (5 AU Myr ') migration rate than previously (Table 3))., We have thus performed simulations with a lower $0.05$ AU $^{-1}$ ) and a higher $5$ AU $^{-1}$ ) migration rate than previously (Table \ref{fastslow}) ).
 Overall. our results are in. good agreement with those of ?.. even for eccentric orbits: the trapping probability increases when the migration rate decreases.," Overall, our results are in good agreement with those of \citet{2003ApJ...598.1321W}, even for eccentric orbits: the trapping probability increases when the migration rate decreases."
 For example. with a low migration rate. a Neptune mass planet traps planetesimals in the 2:1 MMR. while a Saturn mass planet populates the second libration center of the 2:1 resonance.," For example, with a low migration rate, a Neptune mass planet traps planetesimals in the $2$ $1$ MMR, while a Saturn mass planet populates the second libration center of the $2$ $1$ resonance."
 Although the disk resonant shape is modified because the populated MMRs change with the migration rate. the dependence on the mass and eccentricity of the planet remains unchanged: non-resonant planetesimals always dominate at moderate eccentricity and clear resonant patterns are only visible with a planet on a quasi circular orbit.," Although the disk resonant shape is modified because the populated MMRs change with the migration rate, the dependence on the mass and eccentricity of the planet remains unchanged: non-resonant planetesimals always dominate at moderate eccentricity and clear resonant patterns are only visible with a planet on a quasi circular orbit."
 We still note some subtle changes between the different nigration rates., We still note some subtle changes between the different migration rates.
 With a higher migration rate. low mass planets have less time to increase the eccentricities of the non-resonant planetesimals which perturb less the resonant structures.," With a higher migration rate, low mass planets have less time to increase the eccentricities of the non-resonant planetesimals which perturb less the resonant structures."
 With avery low migration rate. planets have more time. during close encounters. to eject non-resonant planetesimals. and resonant planetesimals in the eccentric planet orbit case.," With a very low migration rate, planets have more time, during close encounters, to eject non-resonant planetesimals, and resonant planetesimals in the eccentric planet orbit case."
 Nevertheless. as these changes are small. it is possible to summarize our simulations in the (planet eccentricity: planet mass) parameter space and to discuss the visibility of resonant clumps using only these two parameters. as in Fig. 10..," Nevertheless, as these changes are small, it is possible to summarize our simulations in the (planet eccentricity; planet mass) parameter space and to discuss the visibility of resonant clumps using only these two parameters, as in Fig. \ref{resume}."
 We have assumed in our simulations that the planetesimal disk was initially dynamically cold (ομα= 0.0). but ?.. while using the same hypothesis. pointed out that dynamically warm disks are an interesting alternative to be studied.," We have assumed in our simulations that the planetesimal disk was initially dynamically cold $e_{disk}=0.0$ ), but \citet{2003ApJ...598.1321W}, , while using the same hypothesis, pointed out that dynamically warm disks are an interesting alternative to be studied."
 We have thus extended our work in this direction. assuming disks with," We have thus extended our work in this direction, assuming disks with"
compared with those derived. from. synchrotron spectral ageing arguments.,compared with those derived from synchrotron spectral ageing arguments.
 The latter speeds were estimated for 33 BCR ΤΠ radio sources in the power range 107«ODLO?Wiletse+ by Alexander Leahy (1987) and. Liu (1992). assuming that the source axes lav in the plane of sky. thus providing estimates of the tangential velocities (y= fe) of the lobes.," The latter speeds were estimated for 33 3CRR FRII radio sources in the power range $10^{25}\leq{P_{178}}\leq10^{29}\,{\rm W}{\rm Hz}^{-1}{\rm sr}^{-1}$ by Alexander Leahy (1987) and Liu (1992), assuming that the source axes lay in the plane of sky, thus providing estimates of the tangential velocities $y = v_{\rm t}/c$ ) of the lobes."
 The true mean speed ds related to Ἡ by In the same wav. the standard deviations of the projected and true velocity distributions are related by where yo=(c.," The true mean speed $\overline\beta$ is related to $\overline y$ by In the same way, the standard deviations of the projected and true velocity distributions are related by where $y=v_{\rm t}/c$."
 Phe true mean speed. of the source components derived. from spectral ageing arguments from the above papers is m3j=0.13+ 0.08. compared to a mean value of 3=0.27£0.16 from the observed. clistribution of fractional separation dilferences for the 132 3CRR FRILL sources.," The true mean speed of the source components derived from spectral ageing arguments from the above papers is $\overline{\beta}=0.13\pm0.08$ , compared to a mean value of $\overline{\beta}=0.27\pm0.16$ from the observed distribution of fractional separation differences for the 132 3CRR FRII sources."
 Thus. the mean speed and its standard deviation of the source components are about two times greater than the same quantities derived from spectral ageing arguments.," Thus, the mean speed and its standard deviation of the source components are about two times greater than the same quantities derived from spectral ageing arguments."
 To check this result. the source sample. described. in Section 4 was divided into four equal logarithmic bins in radio power and mean velocity estimates were mace using both techniques (Fig. 1)).," To check this result, the source sample described in Section 4 was divided into four equal logarithmic bins in radio power and mean velocity estimates were made using both techniques (Fig. \ref{fig-1}) )."
 In this analysis. the cosmologies adopted by the above authors have been adopted. {1ο=50 km s| Alpe| and. deceleration parameter gq=0.," In this analysis, the cosmologies adopted by the above authors have been adopted, $H_0=50$ km ${\rm s}^{-1}$ ${\rm Mpc}^{-1}$ and deceleration parameter $q_0=0$."
 There were roughly. equal numbers of sources in each of the bins., There were roughly equal numbers of sources in each of the bins.
 For all four bins. the mean speeds estimated: from. the relativistic svmmetric model are svstematically greater. by [actors of between 1.5 and 5. than the mean speeds obtained from the spectral ageing analvses.," For all four bins, the mean speeds estimated from the relativistic symmetric model are systematically greater, by factors of between 1.5 and 5, than the mean speeds obtained from the spectral ageing analyses."
 The dillerences in the estimates becomes smaller for the most. powerful sources., The differences in the estimates becomes smaller for the most powerful sources.
 Whilst recognising that there are a number of reasons why the spectral ageing estimates might be in error. these discrepancies suggest that the velocities derived. from the svmmetric model are overestimated.," Whilst recognising that there are a number of reasons why the spectral ageing estimates might be in error, these discrepancies suggest that the velocities derived from the symmetric model are overestimated."
 The obvious interpretation of these data is that not all the source asvmmoeltry can be attributed to the relativistic bulk motion of the source components., The obvious interpretation of these data is that not all the source asymmetry can be attributed to the relativistic bulk motion of the source components.
 On the other hand. the diserepaney is less than a factor of two for the most luminous sources and so. although the asvmametries may not be wholly attributed to relativistic elfects. they cannot be neglected.," On the other hand, the discrepancy is less than a factor of two for the most luminous sources and so, although the asymmetries may not be wholly attributed to relativistic effects, they cannot be neglected."
 Laing D(1988) and Carringtone (1988)4 discovered. the important correlation between the degree of depolarisation of a radio lobe and the presence of a one-siclecl racio jet emanating from the active nucleus., Laing (1988) and Garrington (1988) discovered the important correlation between the degree of depolarisation of a radio lobe and the presence of a one-sided radio jet emanating from the active nucleus.
 The sense of the correlation is such that the jet-side lobe is significantly less depolarized than the counterjet-side lobe and can be attributed to the fact that the radiation. from the approaching lobe passes through a smaller path length. of the depolarising halo surrounding the radio source than does the receding lobe., The sense of the correlation is such that the jet-side lobe is significantly less depolarized than the counterjet-side lobe and can be attributed to the fact that the radiation from the approaching lobe passes through a smaller path length of the depolarising halo surrounding the radio source than does the receding lobe.
 The Laine-Carrington cllect thus enables the approaching and receding lobes to be identifed and suggests that the one-sidedness of the jet is the result of relativistic beaming., The Laing-Garrington effect thus enables the approaching and receding lobes to be identifed and suggests that the one-sidedness of the jet is the result of relativistic beaming.
 Phe lobe approaching the observer should be longer than the receding lobe ancl should be the side on which the one-sided radio jet is found., The lobe approaching the observer should be longer than the receding lobe and should be the side on which the one-sided radio jet is found.
 Saikia (1981) used the jet-sice to test. the symmetric model. initially for a small sample of radio galaxies. ancl later for a larger sample of 36 quasars (Saikia. 1984).," Saikia (1981) used the jet-side to test the symmetric model, initially for a small sample of radio galaxies, and later for a larger sample of 36 quasars (Saikia, 1984)."
 For about half the quasars in his sample. the jet was found to Πο on the shorter side.," For about half the quasars in his sample, the jet was found to lie on the shorter side."
 The same effect was also present in the larger sample of quasars and radio galaxies selected. by Scheuer (1995)., The same effect was also present in the larger sample of quasars and radio galaxies selected by Scheuer (1995).
 Phe latest observations of jets from. hieh-resolution VLA observations indicate that. in sources with ouble-sided jets. there is no tendency for the brighter to lie in the longer radio lobe.," The latest observations of jets from high-resolution VLA observations indicate that, in sources with double-sided jets, there is no tendency for the brighter to lie in the longer radio lobe."
 According to the studies listed in Vlection 1. only 18 out of 33 radio galaxies. orA. show 16 expected correlation and. in the quasar sample of Dridle (1994). only 6 out of 13 sources follow the expectation of the relativistic svmmetric model.," According to the studies listed in Section 1, only 18 out of 33 radio galaxies, or, show the expected correlation and, in the quasar sample of Bridle (1994), only 6 out of 13 sources follow the expectation of the relativistic symmetric model."
 The is designed to. take account of the contributions ofboth relativistic and intrinsicenvironmental asvmametries to the structures of, The is designed to take account of the contributions ofboth relativistic and intrinsic/environmental asymmetries to the structures of
"a universal elliciency of GC formation per unit total initial galaxy mass (AleLaugllin1999).. where e is the efficiency. parameter. Af, is (he mass of the visible stellar component of the Sgalaxy. Af...ος is the mass of gas in or around the Sgalactic halo. and VereCi is the mass of the GCS. then Sy must be related with the Iuminositv of the host galaxy. according to: The term LY) accounts for the systematic increase in mass-to-light ratio with galaxy huninositv.","a universal efficiency of GC formation per unit total initial galaxy mass \citep{Mc99}, where $\epsilon$ is the efficiency parameter, $M_*$ is the mass of the visible stellar component of the galaxy, $M_{\rm gas}$ is the mass of gas in or around the galactic halo, and $M_{\rm GC}$ is the mass of the GCS, then $S_N$ must be related with the luminosity of the host galaxy according to: The term $L_{\rm gal}^{0.3}$ accounts for the systematic increase in mass-to-light ratio with galaxy luminosity."
 McLaughlin.(1999) predicts the general trend of ον with luminosity., \citet{Mc99} predicts the general trend of $S_N$ with luminosity.
 The only unknown parameter affecting Say would be the ratio of gas to stellar mass., The only unknown parameter affecting $S_N$ would be the ratio of gas to stellar mass.
 In Chis paper no evidence is found [or a relation between galactic huninositv and ὧν., In this paper no evidence is found for a relation between galactic luminosity and $S_N$.
 On the other hand. i£ a range of possible GC formation efficiencies is allowed. the merger model ean account for the specific [requeney range observed for elliptical galaxies 1992).," On the other hand, if a range of possible GC formation efficiencies is allowed, the merger model can account for the specific frequency range observed for elliptical galaxies \citep{AZ92}."
. I the merger model is to explain the /gh—5. phenomenom. GC formation must be more efficient in the high—ον galaxies.," If the merger model is to explain the $high-S_N$ phenomenom, GC formation must be more efficient in the $high-S_N$ galaxies."
 Forte.Martinez.&Muzzio(1932) and Muzzio(1987) have suggested that in clusters ol galaxies a population of intergalactic GCs (IGCs) should exist., \citet{F82} and \citet{M87} have suggested that in clusters of galaxies a population of intergalactic GCs (IGCs) should exist.
 A galaxy could enlarge its GC population via gravitational capture of IGCs., A galaxy could enlarge its GC population via gravitational capture of IGCs.
 If they. are distributed according to the eravitational potential of the whole galaxy cluster. a relation between 5 and environment is predicted (White1987).," If they are distributed according to the gravitational potential of the whole galaxy cluster, a relation between $S_N$ and environment is predicted \citep{W87}."
. The current dataset shows no indication supporting this prediction., The current data set shows no indication supporting this prediction.
 The fact that no single scenario seems to account for the observed specific [requencies. indicates the historv of each galaxy. should be deduced individually by suitably combining ihe models mentioned above (Woodworth&[Harris2000).," The fact that no single scenario seems to account for the observed specific frequencies, indicates the history of each galaxy should be deduced individually by suitably combining the models mentioned above \citep{WH00}."
. To this aim. it becomes necessary to extend the observational information on each galaxy.," To this aim, it becomes necessary to extend the observational information on each galaxy."
 Besicles (he specific lrequency. ancl (he radial distribution of the GCS. it is necessary {ο know the details of the subpopulations (when these exist) and the kinematics of the GCS.," Besides the specific frequency and the radial distribution of the GCS, it is necessary to know the details of the subpopulations (when these exist) and the kinematics of the GCS."
 In this wax. detailedT observations of the extreme cases mentioned above (NGC 4673 and MCG +5 —31 —063) could be a good starting point because (hese peculiar galaxies can show characteristics found nowhere else. and so offer valuable information in testing the different scenarios.," In this way, detailed observations of the extreme cases mentioned above (NGC 4673 and MCG +5 $-$ 31 $-$ 063) could be a good starting point because these peculiar galaxies can show characteristics found nowhere else, and so offer valuable information in testing the different scenarios."
"the objects in the CDFS have significant detections in X-ray or at The objects are spatially resolved in both J and H, 244m.and, moreover, the J and H band sizes are consistent with each other.","the objects in the CDFS have significant detections in X-ray or at $\mu$ m. The objects are spatially resolved in both J and H, and, moreover, the $J$ and $H$ band sizes are consistent with each other."
" Moreover, it is highly unlikely that all 69 objects are dominated by line emission from active galactic nuclei."," Moreover, it is highly unlikely that all 69 objects are dominated by line emission from active galactic nuclei."
" At least at the present day, low-mass, low-metallicity AGN are exceedingly rare (Izotov&Thuan2008),, much rarer than starbursting dwarf galaxies (Izotovetal.2011)."," At least at the present day, low-mass, low-metallicity AGN are exceedingly rare \citep{izotov08}, much rarer than starbursting dwarf galaxies \citep{izotov11}."
". The implied black hole masses for the objects in our sample, as inferred from their UV continuum luminosities (Shenetal.2008) are ~109Mo at most, when assuming an Eddington accretion of unity."," The implied black hole masses for the objects in our sample, as inferred from their UV continuum luminosities \citep{shen08} are $\sim 10^6~\msol$ at most, when assuming an Eddington accretion of unity."
" At these low masses, at least at the present day, secular processes drive nuclear activity; thus, an unknown accretion mode or triggering mechanism for nuclear activity would have to be invoked in order to explain an extreme change in the relative numbers of AGN and starburst powered emission-line dominated objects."," At these low masses, at least at the present day, secular processes drive nuclear activity; thus, an unknown accretion mode or triggering mechanism for nuclear activity would have to be invoked in order to explain an extreme change in the relative numbers of AGN and starburst powered emission-line dominated objects."
" At these low masses, merging cannot account for this."," At these low masses, merging cannot account for this."
" The starburst hypothesis, on the other hand, places these objects in the realm of dwarf galaxies, and their physical and statistical properties are consistent with the abundances and masses of dwarf galaxies as we will discuss below."," The starburst hypothesis, on the other hand, places these objects in the realm of dwarf galaxies, and their physical and statistical properties are consistent with the abundances and masses of dwarf galaxies as we will discuss below."
 Although nuclear activity cannot be ruled out entirely — and line-strength gradients in star forming z~2 galaxies suggest that weak AGN may contribute to some extent (Trumpetal.2011) — we can safely assume that the observed emission lines are effectively dominated by star formation activity., Although nuclear activity cannot be ruled out entirely – and line-strength gradients in star forming $z\sim 2$ galaxies suggest that weak AGN may contribute to some extent \citep{trump11} – we can safely assume that the observed emission lines are effectively dominated by star formation activity.
model. are populated with galaxies (?)..,"model, are populated with galaxies \citep{Berlind-2002}."
 A fundamental tool for investigating the inlluence of bias on galaxy clustering statistics is the halo occupation clistribution (10D). generally used to determine the average number of galaxies within a dark matter halo as a function of the halo mass.," A fundamental tool for investigating the influence of bias on galaxy clustering statistics is the halo occupation distribution (HOD), generally used to determine the average number of galaxies within a dark matter halo as a function of the halo mass."
 Within the HOD framework. the virialised dark matter haloes with typical overdensities of A~200 (defined with respect to the critical density) are expected to be in approximate dynamical equilibrium and are described in one of three ways: (1) P(N|M). the probability of a halo of given mass M having an average number of galaxies [NO Gi) the relation. between galaxy and dark matter spatial distributions within haloes and (ii) the relation between galaxy and dark matter velocity. distributions within haloes.," Within the HOD framework, the virialised dark matter haloes with typical overdensities of $\Delta \sim 200$ (defined with respect to the critical density) are expected to be in approximate dynamical equilibrium and are described in one of three ways: (i) $|$ M), the probability of a halo of given mass $M$ having an average number of galaxies $\langle N\rangle$; (ii) the relation between galaxy and dark matter spatial distributions within haloes and (iii) the relation between galaxy and dark matter velocity distributions within haloes."
" ""Theoretical studies (e.g. ?????)) based on N-body simulations. hydrodynamic simulations and semi-analvtic models (SAAS) showed that P(N|M) can be well modeled. by a Poissonian distribution in the high halo niass regime (AL,1072 37.) while it. assumes significant sub-Poissonian behavior. which can be modeled by a ""nearest integer (Nint) distribution instead. for lower halo masses."," Theoretical studies (e.g., \citealp{Peacock-2000, Benson-2000a, Berlind-2002, Berlind-2003, Kravtsov-2004}) ) based on N-body simulations, hydrodynamic simulations and semi-analytic models (SAMs) showed that $|$ M) can be well modeled by a Poissonian distribution in the high halo mass regime $M_{\rm h}\gtrsim 10^{12}\ M_{\odot}$ ) while it assumes significant sub-Poissonian behavior, which can be modeled by a “nearest integer (Nint) distribution instead, for lower halo masses."
 These studies also showed (Δον presents a sharp cutolf at low halo mass. a slowly rising plateau and a steep increase of the occupation. number at. high halo mass.," These studies also showed $\langle N\rangle_{\rm M}$ presents a sharp cutoff at low halo mass, a slowly rising plateau and a steep increase of the occupation number at high halo mass."
 X simple way to model the complicated. shape of ANO. avoiding the use of models involving a large number of parameters. is by assuming the existence of two separate galaxy populations within haloes: (i) central. galaxies and (ü) satellite galaxies.," A simple way to model the complicated shape of $\langle N\rangle_{\rm M}$, avoiding the use of models involving a large number of parameters, is by assuming the existence of two separate galaxy populations within haloes: (i) central galaxies and (ii) satellite galaxies."
 This choice is motivated. by reasons wed on hwedrodvnamic simulations (2) and on studies of observed galaxy clusters and groups. which take the xightest cluster galaxies (BCCs) as a dillerent. population rom the rest of the cluster galaxies.," This choice is motivated by reasons based on hydrodynamic simulations \citep{Berlind-2003} and on studies of observed galaxy clusters and groups, which take the brightest cluster galaxies (BCGs) as a different population from the rest of the cluster galaxies."
 Phese two populations can be modeled: separately (e.g... 2273). with the simplest case being modeling the HOD of central galaxies as a step unction. GV.)= above a halo mass limit and a power aw lor satellite galaxies.," These two populations can be modeled separately (e.g., \citealp{Benson-2000a, Kravtsov-2004, Zheng-Coil-2007}) ), with the simplest case being modeling the HOD of central galaxies as a step function $\langle N_{\rm c}\rangle=1$ above a halo mass limit and a power law for satellite galaxies."
1 In this wav the HOD becomes a measure of the combined probability that a halo of mass M rosts à central galaxy and a given number Α of satellite Studies have also been carried. out on the radial. spatial and. velocity clistributions of galaxies and clark matter. the majority of them focused on the latter.," In this way the HOD becomes a measure of the combined probability that a halo of mass $M$ hosts a central galaxy and a given number $N_{\rm s}$ of satellite Studies have also been carried out on the radial, spatial and velocity distributions of galaxies and dark matter, the majority of them focused on the latter."
 For instance. the halo concentration (defined as one of the shape parameters of the radial density profile) and its relation with mass and redshift have been the topic of many investigations in recent vears (ee. 7777)).," For instance, the halo concentration (defined as one of the shape parameters of the radial density profile) and its relation with mass and redshift have been the topic of many investigations in recent years (e.g. \citealp{Nagai-2005, Neto-2007, Gao-2008, Munoz-2011}) )."
 The importance of the concentration resides in its connection with the mean density of the Universe at. the time of collapse. i.e. more concentrated structures formed. earlier in the past. when the Universe was denser (? hereafter NEW. ???)).," The importance of the concentration resides in its connection with the mean density of the Universe at the time of collapse, i.e. more concentrated structures formed earlier in the past, when the Universe was denser \citealp{Navarro-1997} hereafter NFW, \citealp{Neto-2007, Gao-2008, Munoz-2011}) )."
 For this reason. unelerstacing the dependence of concentration on mass ancl redshift and how the relationship between the concentration of dark matter ancl galaxies within halocs evolves with cosmic time. can provide fundamental insight into the formation anc evolution of structures ancl into galaxy formation Several statistical galaxy properties have been used to empirically measure the HOD assuming a power-law form CONIxAL’) for satellite galaxies.," For this reason, understanding the dependence of concentration on mass and redshift and how the relationship between the concentration of dark matter and galaxies within haloes evolves with cosmic time, can provide fundamental insight into the formation and evolution of structures and into galaxy formation Several statistical galaxy properties have been used to empirically measure the HOD assuming a power-law form ${\langle N \rangle}_{\rm M} \propto M^{\beta}$ ) for satellite galaxies."
 For instance. studies of the HOD have been carried out. by using either the luminosity function (e.g.. 2)). the spatial clustering (e.g. 2770) or by counting galaxies within known dark matter haloes. such as rich clusters (e.g... 2?777)).," For instance, studies of the HOD have been carried out by using either the luminosity function (e.g., \citealp{Yang-2008}) ), the spatial clustering (e.g., \citealp{Phleps-2006,Abbas-2010}) ) or by counting galaxies within known dark matter haloes, such as rich clusters (e.g., \citealp{Kochanek-2003, Lin-2004, Collister-Lahav-2005, Popesso-2007, Ho-2009}) )."
 Our study [ocuses on this latter method. which is perhaps the most direct. using a sample of 15 X-ray selected clusters at high redshift.," Our study focuses on this latter method, which is perhaps the most direct, using a sample of 15 X-ray selected clusters at high redshift."
 Unfortunately. there is a. significant disagreement among the results so far obtained ancl firm conclusions are hard to draw.," Unfortunately, there is a significant disagreement among the results so far obtained and firm conclusions are hard to draw."
 For example. the slope of the NoAL relation is scen ranging between 3=0.55 (7). and 3~LY (7): while the average concentration parameter of the galaxy density profile (ο) is found to range between ce2 and CS0OS ?).," For example, the slope of the $N-M$ relation is seen ranging between $\beta=0.55$ \citep{Marinoni-2002} and $\beta \sim 1.7$ \citep{Abbas-2010}; while the average concentration parameter of the galaxy density profile $c_{\rm g}$ ) is found to range between $c_{\rm g}\sim 2$ and $c_{\rm g}\sim 8$ \citep{Carlberg-1997, van-der-Marel-2000, Biviano-2003, Katgert-2004, Lin-2004, Rines-2006, Muzzin-2007a, Biviano-2010}."
 Although the majority. of. studies have been conducted. at. low redshift. these problems of consisteney among dillerent. studies of the HOD. extend to intermediate redshifts (c.g. 222)). therefore there is a continuing motivation to investigate the evolution of the HOD with new galaxy The A-band LE can be used as a surrogate of the galaxy mass function because it is a sensitive probe of the bull xoperties of galaxy populations out to 2c1.5.," Although the majority of studies have been conducted at low redshift, these problems of consistency among different studies of the HOD extend to intermediate redshifts (e.g., \citealp{Lin-2006, Buote-2007, Abbas-2010}) ), therefore there is a continuing motivation to investigate the evolution of the HOD with new galaxy The -band LF can be used as a surrogate of the galaxy mass function because it is a sensitive probe of the bulk properties of galaxy populations out to $z\simeq 1.5$."
 There are several advantages to using the near-infrared light. for such studies: i) A-band (2.2 pini) luminosities broadly reflect he total stellar mass of galaxies. resulting in a ML. ratio hat is insensitive to the star-formation history of earlv- galaxies: ii) stars are easy to remove. as they generally qwe dy1 (Vega). whilst the k-correction makes the observed colours of the great majority of galaxies JAN1 (77): ii) k-correction in the near-infrared bands varies slowly with z and depends weakly on Hubble tvpe (22): iv) he elect of extinction at these wavelengths is significantly smaller than in optical ancl UV. passbands.," There are several advantages to using the near-infrared light for such studies: i) -band $2.2\ {\rm \mu m} $ ) luminosities broadly reflect the total stellar mass of galaxies, resulting in a M/L ratio that is insensitive to the star-formation history of early-type galaxies; ii) stars are easy to remove, as they generally have ${\it J}-{\it K}<1$ (Vega), whilst the k-correction makes the observed colours of the great majority of galaxies ${\it J}-{\it K}>1$ \citep{De-Propris-1999, McCracken-2010}; iii) k-correction in the near-infrared bands varies slowly with $z$ and depends weakly on Hubble type \citep{Poggianti-1997, Bruzual-2003}; iv) the effect of extinction at these wavelengths is significantly smaller than in optical and UV passbands."
 With the adven of infrared surveys like the Survey (UINIDSS. ?)) and Spitzer. it has been possible to carry ou studies of the stellar mass of high-2 galaxies.," With the advent of infrared surveys like the (UKIDSS, \citealp{Lawrence-2007}) ) and Spitzer, it has been possible to carry out studies of the stellar mass of $z$ galaxies."
 However. such μαuclies produced results in contrast with the prediction of SAAIs.," However, such studies produced results in contrast with the prediction of SAMs."
 For instance. 277. and 7. studied the evolution of the cut-off magnitude (A7) of the cluster A-band LE ou to 2 Land found agreement with passive evolving models which have formation redshift (2;~2 5). suggesting tha the bull stellar mass of A cluster galaxies has not increase gsubstantially since z—I.," For instance, \citet{De-Propris-1999, Ellis-2004, Strazzullo-2006} and \citet{Lin-2006} studied the evolution of the cut-off magnitude $K^{\ast}$ ) of the cluster -band LF out to $z\sim 1$ and found agreement with passive evolving models which have formation redshift $z_{\rm f} \sim 2-5$ ), suggesting that the bulk stellar mass of ${\it K}^{\ast}$ cluster galaxies has not increased substantially since $z=1$."
 Several other studies showed tha 1e hieh-mass end of the galaxy mass function seenis to remain pretty much unchanged since z~1 for elliptica ealaxies (c.m. 7. and ?)) and similar results have been obtained for BOCs (?? and. 2)).," Several other studies showed that the high-mass end of the galaxy mass function seems to remain pretty much unchanged since $z\sim1$ for elliptical galaxies (e.g., \citealp{Cimatti-2006} and \citealp{Pozzetti-2010}) ) and similar results have been obtained for BCGs \citealp{Whiley-2008,Collins-2009} and \citealp{Stott-Collins-2010}) )."
" These results suggest a timescale for the mass assemblage of galaxies similar to the age of their component stars. consistent with a monolithic-vpe mocdoel.ancl such activity can be viewed at least as qualitatively as consistent with a 7downsizing"" (7???7? and references therein) process. according to which the more massive earlv-tvpe galaxies end their star formation. anc"," These results suggest a timescale for the mass assemblage of galaxies similar to the age of their component stars, consistent with a monolithic-like model,and such activity can be viewed at least as qualitatively as consistent with a “downsizing” \citealp{Cowie-1996, Thomas-2005, De-Lucia-2007, Stott-2007, Capozzi-2010} and references therein) process, according to which the more massive early-type galaxies end their star formation and"
" These results suggest a timescale for the mass assemblage of galaxies similar to the age of their component stars. consistent with a monolithic-vpe mocdoel.ancl such activity can be viewed at least as qualitatively as consistent with a 7downsizing"" (7???7? and references therein) process. according to which the more massive earlv-tvpe galaxies end their star formation. ancl"," These results suggest a timescale for the mass assemblage of galaxies similar to the age of their component stars, consistent with a monolithic-like model,and such activity can be viewed at least as qualitatively as consistent with a “downsizing” \citealp{Cowie-1996, Thomas-2005, De-Lucia-2007, Stott-2007, Capozzi-2010} and references therein) process, according to which the more massive early-type galaxies end their star formation and"
of the correlation between Ayjjj; and Aycappss 1n regions with varving contributions by the background dust.,"of the correlation between ${\it A_{K,FIR}}$ and ${\it A_{K,2MASS}}$ in regions with varying contributions by the background dust."
 We hope therefore to cisentanele the elect on Aypj; by dust on the far side of the Galaxy from other factors that may. Lead to discrepancies in the Agpijes. Aycapiss relation.," We hope therefore to disentangle the effect on ${\it A_{K,FIR}}$ by dust on the far side of the Galaxy from other factors that may lead to discrepancies in the ${\it A_{K,FIR}}~vs.$ ${\it A_{K,2MASS}}$ relation."
 Fig., Fig.
 ll shows the Anier: os. ντος relation in the region 5>[bfc3.," 11 shows the ${\it A_{K,FIR}}$ $vs.$ ${\it A_{K,2MASS}}$ relation in the region $5^{\circ}>|{\it b}|>3^{\circ}$."
 The upper panels correspond. to cells in the northern Galactic strip ([(|< 57. 5°>b 37). whereas the lower panels correspond. to the southern (Πέ57. 3mb 5°) strip.," The upper panels correspond to cells in the northern Galactic strip $|\ell|<5^{\circ}$ , $5^{\circ}>{\it b}>3^{\circ}$ ), whereas the lower panels correspond to the southern $|\ell|<5^{\circ}$, $-3^{\circ}>{\it b}>-5^{\circ}$ ) strip."
 The νοτωντνΓι ratio histograms in panels (11a) and (lle) show that the relative extinction values have peaks at 0.75 (75 )) and 0.72 ()) in the northern and. southern strips. respectively.," The $A_{K,2MASS}/ A_{K,FIR}$ ratio histograms in panels (11a) and (11c) show that the relative extinction values have peaks at 0.75 (75 ) and 0.72 (72 ) in the northern and southern strips, respectively."
 These well-defined. peaks imply a good linear. correlation )etween the two extinction values. as confirmed. in. panels (11b) and (11d).," These well-defined peaks imply a good linear correlation between the two extinction values, as confirmed in panels (11b) and (11d)."
 For the northern and. southern strips we jwe the correlations Ayjig=1.33ARayes ancl ALApy=1901. pass. respectively. where the angular cocLlicients were derived from the median «κοτοςνριιῃ ratio values.," For the northern and southern strips we have the correlations ${\it A_{K,FIR}}=1.33 {\it A_{K,2MASS}}$ and ${\it A_{K,FIR}}=1.39 {\it A_{K,2MASS}}$ , respectively, where the angular coefficients were derived from the median $A_{K,2MASS}/ A_{K,FIR}$ ratio values."
 In this region. the mocel predictions indicate that. for both northern and southern strips. the ratio of extinction on foreground of the Galactic to total extinction should. be96'4.. independent of the Sun's displacement from the cise mid-plane.," In this region, the model predictions indicate that, for both northern and southern strips, the ratio of extinction on foreground of the Galactic to total extinction should be, independent of the Sun's displacement from the disc mid-plane."
 From the median elysassclpig values we infer that a typical value for this ratio is z το...," From the median $A_{K,2MASS}/ A_{K,FIR}$ values we infer that a typical value for this ratio is $\approx$ ."
 Fhorefore this discrepancy between model and observed. extinction ratio suggests that the contribution [rom background. dust cannot explain the observed zlouassApig ratio.," Therefore this discrepancy between model and observed extinction ratio suggests that the contribution from background dust cannot explain the observed $A_{K,2MASS}/ A_{K,FIR}$ ratio."
 Arce Goodman (1999) also found a linear relation between the dust. emission. extinction ancl that. derived from the stellar content in the Taurus Dark Cloud. with a slope that varies in the range 10-10.," Arce Goodman (1999) also found a linear relation between the dust emission extinction and that derived from the stellar content in the Taurus Dark Cloud, with a slope that varies in the range 1.3-1.5."
 They attributed this dillerence to the dust column censitv es. recddening calibration [from Schlegel ct al. (, They attributed this difference to the dust column density $vs.$ reddening calibration from Schlegel et al. (
1998). which may vield an overestimated Ay 0.5.,"1998), which may yield an overestimated ${\it A_V}>$ 0.5."
 Note that the slopes in our two strips (with 2 >chyc 10) are similar to those of Arce Goodman (1999). whichcorroborates the idea of a zlrig calibration elfect.," Note that the slopes in our two strips (with 2 $> A_V >$ 10) are similar to those of Arce Goodman (1999), whichcorroborates the idea of a $A_{K,FIR}$ calibration effect."
 We point out that assuming a significantIy lower value [or Ay would compress the elipij scale. decreasing the dillerences. between «κει and clayeaass dn Fie.," We point out that assuming a significantly lower value for $R_V$ would compress the $A_{K,FIR}$ scale, decreasing the differences between $A_{K,FIR}$ and $A_{K,2MASS}$ in Fig."
 11., 11.
 Llowever. observational constraints on Z: do not. support this possibility.," However, observational constraints on $R_V$ do not support this possibility."
 Indeed. Could et al. (," Indeed, Gould et al. ("
2001) estimated. Ly = AlySEV1) = 2.4 (Rym 3.0). from VLA. colours of 146 Daade's Window CG and |x giants.,"2001) estimated $R_{VI}$ = $A_V /E(V-I)$ = 2.4 $R_V \approx$ 3.0), from $VIK$ colours of 146 Baade's Window G and K giants."
" Stanck (1996) obtained 0, = 2.50.1 (I8 3.1). using Baacle’s Window red clump giants."," Stanek (1996) obtained $R_{VI}$ = $\pm 0.1$ $R_V = 3.1$ ), using Baade's Window red clump giants."
 For stellar fields and reddened metal-rich elobular clusters throughout the Bulge. values of Ry=3.53.6 have been emploved. (Lerncrup 1988. Barbuy οἱ al.," For stellar fields and reddened metal-rich globular clusters throughout the Bulge, values of $R_V = 3.5-3.6$ have been employed (Terndrup 1988, Barbuy et al."
 1998)., 1998).
 We conclude that typical Z5: values in. Bulec directions cannot be significantly lower than £:=3.1., We conclude that typical $R_V$ values in Bulge directions cannot be significantly lower than $R_V = 3.1$.
 Assuming that the discrepancy between model and (υπονο! extinction ratio in these strips is entirely due to a calibration. problem. we estimate a calibration factor of to be applied to the slicpg values.," Assuming that the discrepancy between model and observed extinction ratio in these strips is entirely due to a calibration problem, we estimate a calibration factor of to be applied to the $A_{K,FIR}$ values."
 Panels (11b) and (11d) also reveal à strong asymmetry between the northern and southern strips., Panels (11b) and (11d) also reveal a strong asymmetry between the northern and southern strips.
 values reach up to —1.0 in the north. with most cells in the range 0.2<οντως0.7.," values reach up to =1.0 in the north, with most cells in the range $0.2< {\it A_{K,2MASS}}<0.7$."
 In the south. there are few cells with Aycoabdss0.4.," In the south, there are few cells with ${\it A_{K,2MASS}}>0.4$."
 In Sect. 4.4. we discuss the north-south asymmetry in more cetail.," In Sect 4.4, we discuss the north-south asymmetry in more detail."
 As we consider regions closer to the Galactic plane. several predictions can be made regarding the extinction. values derived. from the 2ALASS and. DIRBE/IRAS data.," As we consider regions closer to the Galactic plane, several predictions can be made regarding the extinction values derived from the 2MASS and DIRBE/IRAS data."
 First. one obviously expects a general increase in both iMosiiss and κιν values.," First, one obviously expects a general increase in both ${\it A_{K,2MASS}}$ and ${\it A_{K,FIR}}$ values."
 Further. the relation between the two should. increasingly depart from the identity line. as the contribution of background. dust. is enhanced.," Further, the relation between the two should increasingly depart from the identity line, as the contribution of background dust is enhanced."
 Finally. line-of-sight variations should also grow in amplitude. as more individual dust clouds. with variable dust. densities and temperatures. are expected. to lie along low latitude clirections.," Finally, line-of-sight variations should also grow in amplitude, as more individual dust clouds, with variable dust densities and temperatures, are expected to lie along low latitude directions."
 All these predictions. are. confirmed. by inspection of rie.12. which shows the comparison between JAoirissand Ancry. lor the region 3°26z I.," All these predictions are confirmed by inspection of Fig.12, which shows the comparison between ${\it A_{K,2MASS}}$and ${\it A_{K,FIR}}$ for the region $3^{\circ}>{\it b}>1^{\circ}$ ."
 The Ayου. histograms in panels (12a) and (120) are dominated: by double peak distributions. at lyourassdpig = 0.75 and 0.60 for the northern strip. and at olsrτοςτει =," The $A_{K,2MASS}/A_{K,FIR}$ histograms in panels (12a) and (12c) are dominated by double peak distributions, at $A_{K,2MASS}/A_{K,FIR}$ = 0.75 and 0.60 for the northern strip, and at $A_{K,2MASS}/A_{K,FIR}$ ="
For galaxies. black hole mass also correlates with buele luminosity.,"For galaxies, black hole mass also correlates with bugle luminosity."
 Such a correlation is expected vot to present iu elobular clusters., Such a correlation is expected yet to present in globular clusters.
 The huninosity of elobular clusters is plotted against black hole mass in Fie., The luminosity of globular clusters is plotted against black hole mass in Fig.
 2., 2.
 The solid line gives the best fit of galaxies by WG2001., The solid line gives the best fit of galaxies by KG2001.
 The elobular clusters fall around tle solid hue with a considerable scatter., The globular clusters fall around the solid line with a considerable scatter.
 This cousisteucy reinforces the possibility that the globular clusters harbor black holes aud satisfy the same Mpg — 0 correlation as that of ealaxies., This consistency reinforces the possibility that the globular clusters harbor black holes and satisfy the same $_{BH}$ – $\sigma$ correlation as that of galaxies.
 Tt isa pleasure to thank Prof. N.-P. Wu for lis creative advice aud discussions., It is a pleasure to thank Prof. X.-P. Wu for his creative advice and discussions.
 The author is extremely erateful to Prof. L.-C. Dene for excellent sugecstions aud so much help., The author is extremely grateful to Prof. L.-C. Deng for excellent suggestions and so much help.
 Maux thauks to the staffs of stars and stellar svstenis eroup for their very kind help., Many thanks to the staffs of stars and stellar systems group for their very kind help.
 The author also thank the referee Prof. Z.-CG. Dene for helpful suggestions that improve the paper., The author also thank the referee Prof. Z.-G. Deng for helpful suggestions that improve the paper.
 This work is supported in part by the Ministry of Science aud Technology of China through eraut. 6119990751., This work is supported in part by the Ministry of Science and Technology of China through grant G19990754.
Many such models break the population down into contributions from weakly or non-evolving normal galaxies. together with a population of starburst galaxies subject to strong evolution in space density and luminosity.,"Many such models break the population down into contributions from weakly or non-evolving normal galaxies, together with a population of starburst galaxies subject to strong evolution in space density and luminosity."
 The model of Franceshini et al. (, The model of Franceshini et al. (
2001) is one such model. which takes as its starting point a local luminosity function.,"2001) is one such model, which takes as its starting point a local luminosity function."
 However. when translated to with the assumed infrared SEDs. the resulting breakdown into normal and starburst galaxies (Fig.," However, when translated to with the assumed infrared SEDs, the resulting breakdown into normal and starburst galaxies (Fig."
 || of Franceschini et al., 11 of Franceschini et al.
 2001) is completely different from the one we assume. with the starbursts comprising a decreasing fraction of the whole for Lov10Η...," 2001) is completely different from the one we assume, with the starbursts comprising a decreasing fraction of the whole for $L_{\rm{60}} > 10^{10}$."
 When translated to (Gruppioni et al., When translated to (Gruppioni et al.
 2005). the Franceschini et al.," 2005), the Franceschini et al."
 model is at variance| with the redshift breakdown ofSpitzer source counts reported by Le Floc’h et al. (, model is at variance with the redshift breakdown of source counts reported by Le Floc'h et al. (
2005).,2005).
 The Lagache et al. (, The Lagache et al. (
2004) model decomposition of the local Iuminosity function into normal and starburst galaxies is similar to ours. but again there are discrepancies with respect to the luminosity functions reported by Le Floc'h et al. (,"2004) model decomposition of the local luminosity function into normal and starburst galaxies is similar to ours, but again there are discrepancies with respect to the luminosity functions reported by Le Floc'h et al. ("
2005).,2005).
 In the absence of a definitive alternative model and to avoid the uncertainties arising from K-corrections in the mid-infrared. we prefer to base our predictions starting from the local 1.4 GHz luminosity function plus pure luminosity evolution for both sub-populations. as already discussed.," In the absence of a definitive alternative model and to avoid the uncertainties arising from K-corrections in the mid-infrared, we prefer to base our predictions starting from the local 1.4 GHz luminosity function plus pure luminosity evolution for both sub-populations, as already discussed."
 The simulation was performed by assuming no further evolution in the luminosity function beyond redshift ;=1.5., The simulation was performed by assuming no further evolution in the luminosity function beyond redshift $z=1.5$.
 Taken at face value this is clearly unrealistic. but it gives the user freedom to impose their own form of high redshift decline by selectively. filtering the catalogue during post-processing.," Taken at face value this is clearly unrealistic, but it gives the user freedom to impose their own form of high redshift decline by selectively filtering the catalogue during post-processing."
 At the present time. observations show that the star formation rate density is essentially constant from 2~1.5 to z~+ and then to drops off sharply above this redshift (see e.g. Hopkins Beacom 2006).," At the present time, observations show that the star formation rate density is essentially constant from $z \sim 1.5$ to $z \sim 4$ and then to drops off sharply above this redshift (see e.g. Hopkins Beacom 2006)."
" As our default post-processing option. we model this high-redshift fall-off using the piece-wise power-law model of Hopkins Beacom. in which the star-formation rate density falls off as (112)"""" above :=4.8, e The interplay between bursts of star formation and black hole accretion is of considerable interest for studies of galaxy evolution. and one aspect of this concerns hybrid galaxies in which both processes make significant contributions to the observed radio emission."," As our default post-processing option, we model this high-redshift fall-off using the piece-wise power-law model of Hopkins Beacom, in which the star-formation rate density falls off as $(1+z)^{-7.9}$ above $z=4.8$ $\bullet$ The interplay between bursts of star formation and black hole accretion is of considerable interest for studies of galaxy evolution, and one aspect of this concerns hybrid galaxies in which both processes make significant contributions to the observed radio emission."
 Unfortunately. such objects cannot be isolated in the simulations due to the use of separate luminosity functions for the AGN and the starbursts. but they are almost certainly accounted for by virtue of the way in which the luminosity functions were constructed.," Unfortunately, such objects cannot be isolated in the simulations due to the use of separate luminosity functions for the AGN and the starbursts, but they are almost certainly accounted for by virtue of the way in which the luminosity functions were constructed."
 A fraction of the luminous starbursts are likely to be mis-classified obscured AGN tor at least have a significant AGN contribution). and similarly some of the radio emission in the low uminosity AGN (especially the radio-quiet population) is likely to be related to star formation.," A fraction of the luminous starbursts are likely to be mis-classified obscured AGN (or at least have a significant AGN contribution), and similarly some of the radio emission in the low luminosity AGN (especially the radio-quiet population) is likely to be related to star formation."
 The most likely upshot of this is that. in the simulation as a whole. the latter populations may have been double counted to a certain extent. thereby slightly overproducing he source counts.," The most likely upshot of this is that, in the simulation as a whole, the latter populations may have been double counted to a certain extent, thereby slightly overproducing the source counts."
 The semi-analytical simulations. in contrast. will include hybrid galaxies.," The semi-analytical simulations, in contrast, will include hybrid galaxies."
 A further instance of potential double counting concerns the fact that the hard X-ray luminosity unction used for the radio-quiet AGN implicitly includes some contribution from the radio-loud AGN. so that the latter are in effect counted twice.," A further instance of potential double counting concerns the fact that the hard X-ray luminosity function used for the radio-quiet AGN implicitly includes some contribution from the radio-loud AGN, so that the latter are in effect counted twice."
 However. this effect is likely dwarfed by he uncertainty in the Compton-thick correction factor applied to he X-ray luminosity function and we therefore make no further allowance for it.," However, this effect is likely dwarfed by the uncertainty in the Compton-thick correction factor applied to the X-ray luminosity function and we therefore make no further allowance for it."
 As described in section 2.4. radio-loud AGN are initially drawn from a 151 MHz luminosity function for steep-spectrum. lobe-dominated sources.," As described in section 2.4, radio-loud AGN are initially drawn from a 151 MHz luminosity function for steep-spectrum, lobe-dominated sources."
 We use this as the input parent population for an orientation-based unification model which enables us to assign source structures and radio spectra in a physically-motivated manner., We use this as the input parent population for an orientation-based unification model which enables us to assign source structures and radio spectra in a physically-motivated manner.
 Some of this procedure follows the prescription of Jackson Wall (1999)., Some of this procedure follows the prescription of Jackson Wall (1999).
" The steps in our process are as follows: (i) Sources are assigned a true linear size. {δινω drawn at random from a uniform distribution [O.Do(1.|2)ET]. where LX,=1 Mpe: this assumes that the sources expand with uniform velocity until they reach a size equal to the redshift-dependent upper envelope of the projected linear size distribution measured by Blundell. Rawlings Willott (1999). ("," The steps in our process are as follows: (i) Sources are assigned a true linear size, $D_{\rm{true}}$ , drawn at random from a uniform distribution $D_{\rm{0}} (1+z)^{-1.4}$ ], where $D_{\rm{0}}=1$ Mpc; this assumes that the sources expand with uniform velocity until they reach a size equal to the redshift-dependent upper envelope of the projected linear size distribution measured by Blundell, Rawlings Willott (1999). ("
i) The angle between the jet axis and the observer's is drawn from a uniform distribution in cos 8. and the jet axis is given a random position angle on the sky. (,"ii) The angle between the jet axis and the observer's line-of-sight is drawn from a uniform distribution in cos $\theta$, and the jet axis is given a random position angle on the sky. ("
iil) The ratio of the intrinsic core to extended luminosities. defined at 1.4 GHz in the rest-frame. is given by Pop=10”. where .r is drawn from a Gaussian distribution of mean μοι and ao=0.5. truncated at abs(.r)10 to avoid numerical problems.,"iii) The ratio of the intrinsic core to extended luminosities, defined at 1.4 GHz in the rest-frame, is given by $R_{\rm{CL}}=10^x$, where $x$ is drawn from a Gaussian distribution of mean $x_{\rm{med}}$ and $\sigma=0.5$, truncated at $(x)>10$ to avoid numerical problems."
 The numerical values of μμ we use are given at the end of this (iv) A relativistic beaming model is used to derive the observed core:extended. flux ratio. Pops=ReL(A). where D(8)—σαdeos&)7|(11deos()7:4 =νί21)/ and ~ is the Lorentz factor of the jet. (," The numerical values of $x_{\rm{med}}$ we use are given at the end of this (iv) A relativistic beaming model is used to derive the observed core:extended flux ratio, $R_{\rm{OBS}}= R_{\rm{CL}} B(\theta)$, where $B(\theta) = \frac{1}{2}[(1 - \beta cos \theta)^{-2} + 
(1 + \beta cos \theta)^{-2}]$; $\beta= \sqrt(\gamma^2 - 1)/\gamma$ and $\gamma$ is the Lorentz factor of the jet. ("
v) The extended emission is modelled with a power-law spectrum £x7 while the core spectrum is modelled with some curvature: log£5.ay|OudogeaotlosgUY.,"v) The extended emission is modelled with a power-law spectrum $F_{\rm{\nu}} \propto \nu^{-0.75}$, while the core spectrum is modelled with some curvature: $\log F_{\rm{\nu}} = a_{\rm{c0}} + a_{\rm{c1}}\log \nu + a_{\rm{c2}} (\log \nu)^{2}$."
" For vin GHz. a,=0.01. @2=—0.29. as measured by Jarvis Rawlings (2000) from a sample of flat spectrum quasars (au sets the normalization)."," For $\nu$ in GHz, $a_{\rm{c1}}=0.07$, $a_{\rm{c2}}=-0.29$, as measured by Jarvis Rawlings (2000) from a sample of flat spectrum quasars $a_{\rm{c0}}$ sets the normalization)."
 The observational dati on which these fits were based typically extend up to 10 or 20 GHz. implying that the model SEDs cannot be simply extrapolated to higher frequencies (e.g. to the WMAP bands above 20 GHz). (," The observational data on which these fits were based typically extend up to 10 or 20 GHz, implying that the model SEDs cannot be simply extrapolated to higher frequencies (e.g. to the WMAP bands above 20 GHz). ("
vi) For FRIs. the extended flux distribution on the sky is modelled as two coaxial elliptical lobes of uniform surface brightness. extending from the point source core. and each with a major axis length equal to half the projected linear size.,"vi) For FRIs, the extended flux distribution on the sky is modelled as two coaxial elliptical lobes of uniform surface brightness, extending from the point source core, and each with a major axis length equal to half the projected linear size."
 The axial ratio of the lobes is drawn from a uniform distribution [0.2.1]. (," The axial ratio of the lobes is drawn from a uniform distribution [0.2,1]. ("
vii) For FRIIs. the inner edges of the lobes are offset from the core by a distance fo.PLS. where LS is the projected linear size of the whole source and f. is drawn from the uniform distribution [0.2.0.8].,"vii) For FRIIs, the inner edges of the lobes are offset from the core by a distance $f\times PLS$, where $PLS$ is the projected linear size of the whole source and $f$ is drawn from the uniform distribution [0.2,0.8]."
 A fraction fuz of the extended flux in the FRIIs is assigned to point-source hotspots positioned at the ends of the lobes. where inaccordance with a correlation found by Jenkins McEllin (1977). (the seatter is modelled with a uniform distribution).," A fraction $f_{\rm{HS}}$ of the extended flux in the FRIIs is assigned to point-source hotspots positioned at the ends of the lobes, where inaccordance with a correlation found by Jenkins McEllin (1977) (the scatter is modelled with a uniform distribution)."
 The hotspots are assumed to have the same, The hotspots are assumed to have the same
"whereLj,=Ly and L,=Lp.",where$L_1=L_{\rm A}$ and $L_4=L_{\rm B}$.
 We can calculate thevalues of5. τοι and 554 from equations (2)) and. (3)) alter the parameters of shells are given.," We can calculate thevalues of$\gamma$, $\gamma_{21}$ and $\gamma_{34}$ from equations \ref{relgam}) ) and \ref{dyne}) ) after the parameters of shells are given."
 In four limit cases. these equations can be solved. analytically (Yu Dai 2009).," In four limit cases, these equations can be solved analytically (Yu Dai 2009)."
" For 2a 09 d£ LaufLam(GA)Gaf£54). we have ce,=safeX losalecl/Tfsiande —540—23€) which means the forward shock is relativistic and the reverse shock is Newtonian: (2) i£ 16«Ljí4/L,(1/16)(54SEN we can obtain. 5»,=fi)EM1/224071/2.m lossDpoptLn1/2 and 5= so both the two shocks are relativistic: (3) if Lifly«xY. we ect ro,lm—¢LTS: sap0m54/25, ands=,(1|28). so the forward shock is Newtonian and the reverse shock is relativistic."," For $\gamma_4\gg\gamma_1$, (1) if ${L_4/ L_1}\gg{(1/7)}\left({\gamma_4/\gamma_1}\right)^4$, we have $\gamma_{21}={\gamma_4/2\gamma_1}\gg1$ , $\gamma_{34}-1\approx{\gamma_4^2/7f\gamma_1^2}$ and $\gamma=\gamma_4(1-\sqrt{2\xi})$, which means the forward shock is relativistic and the reverse shock is Newtonian; (2) if $16\ll{L_4/
L_1}\ll{(1/16)}\left({\gamma_4/\gamma_1}\right)^4$, we can obtain $\gamma_{21}={f^{1/4}\gamma_4^{1/2}/2\gamma_1^{1/2}}\gg1$ , $\gamma_{34}={\gamma_4^{1/2}/2f^{1/4}\gamma_1^{1/2}}\gg1$ and $\gamma=f^{1/4}\gamma_1^{1/2}\gamma_4^{1/2}$, so both the two shocks are relativistic; (3) if ${L_4/ L_1}\ll7$, we get $\gamma_{21}-1\approx{f\gamma_4^2/7\gamma_1^2}=\xi$, $\gamma_{34}={\gamma_4/2\gamma_1}$ and $\gamma=\gamma_1(1+\sqrt{2\xi})$, so the forward shock is Newtonian and the reverse shock is relativistic."
 Finally. (4) for 547 54. both the two shocks are Newtonian.," Finally, (4) for $\gamma_4\approx\gamma_1$ , both the two shocks are Newtonian."
" Since 1. 74. and f are unchanged with the moving of the shells. the values of = so, and 554 are constant before the shocks eross the shells (Yu Dai 2009)"," Since $\gamma_1$, $\gamma_4$, and $f$ are unchanged with the moving of the shells, the values of $\gamma$ $\gamma_{21}$ and $\gamma_{34}$ are constant before the shocks cross the shells (Yu Dai 2009)."
" Following Dai Lu (2002). the total number of the electrons swept-up by the forward ancl reverse shocks during a period of df can be expressed by Nos=2V9£LOI(aπου” and NN,=Lpol(tsumpe}. respectively (Yu. Wang Dai 2009)"," Following Dai Lu (2002), the total number of the electrons swept-up by the forward and reverse shocks during a period of $\delta t$ can be expressed by $N_{e,2}={2\sqrt{2\xi} L_A\delta
t/\left(\psi(z)\gamma_1m_pc^2\right)}$ and $N_{e,3}={L_B\delta
t/\left(\psi(z)\gamma_4m_pc^2\right)}$, respectively (Yu, Wang Dai 2009)."
 The forward ancl reverse shocks can accelerate particles to high energies., The forward and reverse shocks can accelerate particles to high energies.
 Following Sari et al. (, Following Sari et al. (
L998). we assume that the energies of the hot electrons and magnetic fields are [ractions ο and cg of the total internal energy. respectively.,"1998), we assume that the energies of the hot electrons and magnetic fields are fractions $\epsilon_e$ and $\epsilon_B$ of the total internal energy, respectively."
 Thus. the strength of the magnetic fields. is ∠⇂↕≻⇂↓⋅↕∣⋡⇂⇂∣↕⋖⋟↓↕∪⇂⋅↥↓∐⋅⋡∖↓↥⋖⋈∙↓∡−⋯∙≼∙⋖⊾↓⋖⋅↓," Thus, the strength of the magnetic fields is $B_i=\left(8\pi
\epsilon_{B,i}e_i\right)^{1/2}, ~~i=2,3$."
"⋅⋜⊔⋯⇂⋖⊾⇂⋖⊾≼∙∣↓⋅∪⊔⋡∖⋡↙∣⊔↴↙∣⋎↴∖..ο;−⋅ We assume a power-law 5."" for sa.Zaya, (Sari et al."," We assume a power-law distribution of the shock-accelerated electrons, $dn_e/d\gamma_e\propto\gamma_e^{-p}$ for $\gamma_e\geq\gamma_{e,m}$ (Sari et al."
 1998)., 1998).
 The random. Lorentz factor of electrons in regions 2 or 3 is determined by SeanT(QUEEStoytpoLLaitB(PL). where DH equals to 521 or 534.," The random Lorentz factor of electrons in regions 2 or 3 is determined by $\gamma_{e,m,i}=\epsilon_{e,i}{m_p\over
m_e}{(p-2)\over(p-1)}(\Gamma-1)$, where $\Gamma$ equals to $\gamma_{21}$ or $\gamma_{34}$."
" 1n both shockedne, regions. the hot electrons. with energies above 5,ο lose most of their energies during a cooling time of. where the cooling Lorentz [actor is determined bv οίπιοι)(σι; ot)."," In both shocked regions, the hot electrons with energies above $\gamma_{e,c,i}m_ec^2$ lose most of their energies during a cooling time $\delta t$ , where the cooling Lorentz factor is determined by ${\gamma}_{e,c,i}={6\pi
m_ec\psi(z)/\left(\sigma_T{B}_i^2\gamma \delta t\right)}$ ."
 The two characteristic frequencies ancl a peak flux density are (Sari et al., The two characteristic frequencies and a peak flux density are (Sari et al.
 1998: Wijers Galama 1999) where dp=c(l|2)/Llo[|————— is the Iuminosity. distance of the source and fp) is a function of p. dor p=2.2. &(p)=0.6 (Wijers Galama 1999).," 1998; Wijers Galama 1999) where $d_L=c(1+z)/H_0\int_0^z\frac{dz'}{\sqrt{\Omega_M(1+z')^3+\Omega_\Lambda}}$ is the luminosity distance of the source and $\Phi(p)$ is a function of $p$, for $p=2.2$, $\Phi(p)\approx 0.6$ (Wijers Galama 1999)."
" In the calculation. we use Qa,=0.3. O4=0.7 and 042-10 km +."," In the calculation, we use $\Omega_M=0.3$, $\Omega_\Lambda=0.7$ and $H_0$ =70 km $^{-1}$ $^{-1}$."
 gq. is the electron charge anc op is the Thomson+ cross section., $q_e$ is the electron charge and $\sigma_T$ is the Thomson cross section.
 The synchrotron spectrum can be written as (Sari et al., The synchrotron spectrum can be written as (Sari et al.
 1998) where £j=ΑΗνι). vy=max(tui.1). and q—2 lor £a«uus and qo plor οD uu.," 1998) where $\nu_{l}=\min(\nu_{m,i},\nu_{c,i})$, $\nu_{h}=\max(\nu_{m,i},\nu_{c,i})$, and $q=2$ for $\nu_{c,i}<\nu_{m,i}$ and $q=p$ for $\nu_{c,i}>\nu_{m,i}$ ."
 There aretwo peaks in the spectrum of Swift J1644|57. far-infrared (PER) and hard X-ray. peaks.," There aretwo peaks in the spectrum of Swift J1644+57, far-infrared (FIR) and hard X-ray peaks."
 In order to fit the spectrum. we focus on the case (3) of internal shock model in section 2.1. in which the reverse shock is relativistic and the forward shock is Newtonian.," In order to fit the spectrum, we focus on the case (3) of internal shock model in section 2.1, in which the reverse shock is relativistic and the forward shock is Newtonian."
" In the rest of the paper we denote Q=10,", In the rest of the paper we denote $Q=10^xQ_x$.
" For illustration purpose. we set L,—Lj,L=10""ergs to σι= 1000. 5,= I0. QuomQu=o0.5 and ep»=cgacg 0.1.Aceorcing to observation. we usedf—100 s. the variability. timescale of flare (Bloom etal."," For illustration purpose, we set $L_4=L_1=L=10^{47.0}\rm erg~s^{-1}$ , $\gamma_4=1000$ , $\gamma_1=10$ , $\epsilon_{e,2}=\epsilon_{e,3}=\epsilon_{e}=0.5$ and $\epsilon_{B,2}=\epsilon_{B,3}=\epsilon_{B}=0.1$ .According to observation, we use$\delta t\sim 100$ s, the variability timescale of flare (Bloom etal."
 2011: Burrows et al., 2011; Burrows et al.
 2011)., 2011).
 The collision radius is A254~251094feos) lO0Hem. which is consistent with the X-rayemission paclius determinedfrom observation (Bloom οἱ al.," The collision radius is $R_{\rm col}\sim 2\gamma_1^2c\delta t/\psi(z)\sim5\times
10^{14}$ cm, which is consistent with the X-rayemission radius determinedfrom observation (Bloom et al."
 2011)., 2011).
 The Lorentz factor of merged shellis ~ 14., The Lorentz factor of merged shellis $\gamma\sim 14$ .
 Using equation (4)). we can obtain the following expressions for the reverse shock," Using equation \ref{vmc}) ), we can obtain the following expressions for the reverse shock"
of searching for either of these forms of activity is to utilise the radio emission.,of searching for either of these forms of activity is to utilise the radio emission.
 AGN are extremely luminous at radio wavelengths. with luminosities reaching as hieh as about 1077 HEIz.to at GCllIz.," AGN are extremely luminous at radio wavelengths, with luminosities reaching as high as about $10^{28}$ $^{-1}$ at GHz."
 Normal (ie non-AGN) galaxies have a range of observed radioluminosities. »etween about 1077 and 1075 fat ο with starbursting galaxies filling the upper half of that range.," Normal (ie non-AGN) galaxies have a range of observed radioluminosities, between about $10^{18}$ and $10^{24}$ $^{-1}$ at GHz, with starbursting galaxies filling the upper half of that range."
 The radio emission arising from these normal galaxies comes edominantlvy from the svnachrotron emission of particles accelerated. in. supernova shocks (e.g. Condon 1992. and references therein). with a smaller. thermal contribution rom regions. ancl so essentially rellects the current star ormation of the galaxy.," The radio emission arising from these normal galaxies comes predominantly from the synchrotron emission of particles accelerated in supernova shocks (e.g. Condon 1992 and references therein), with a smaller thermal contribution from regions, and so essentially reflects the current star formation of the galaxy."
 At redshifts 2=0.2 the radio luminosities of luminous starburst galaxies ancl weak AGN correspond to microjansky lux density levels. and over recent vears numerous {ην evel radio surveys of the field have been carried. out. (?:7T:2:17:08 ?)..," At redshifts $z \gta 0.2$ the radio luminosities of luminous starburst galaxies and weak AGN correspond to microjansky flux density levels, and over recent years numerous $\mu$ Jy level radio surveys of the field have been carried out \cite{don87,fom91,win93,ric98,mux99,gar00a}."
 The integral radio source count shows an upturn below a couple of mv. indicating the emergence of a new population of radio sources at microjansky levels (?:?): indeed. à laree proportion of the srJv radio source »opulation have been shown to be associated with starburst galaxies and normal spiral galaxies at. substantial recdshifts >0.2 (e.g. Windhorst et al 1995).," The integral radio source count shows an upturn below a couple of mJy, indicating the emergence of a new population of radio sources at microjansky levels \cite{win84,win85}; indeed, a large proportion of the $\mu$ Jy radio source population have been shown to be associated with starburst galaxies and normal spiral galaxies at substantial redshifts $z
\gg 0.2$ (e.g. Windhorst et al \nocite{win95}."
. The steep slope of he differential source counts of μον radio sources. similar o that of the very powerful. radio galaxies ancl quasars. implies that this population of sources. if it is a sinele population. is undergoing strong cosmological evolution with either the density or luminosity of the sources being higher in the past.," The steep slope of the differential source counts of $\mu$ Jy radio sources, similar to that of the very powerful radio galaxies and quasars, implies that this population of sources, if it is a single population, is undergoing strong cosmological evolution with either the density or luminosity of the sources being higher in the past."
 This is in broad agreement with determinations of the evolution of the cosmic star formation rate (?).. suggesting that studies of μον radio sources are important for understanding the star formation history and evolution of ordinary galaxies (?)..," This is in broad agreement with determinations of the evolution of the cosmic star formation rate \cite{mad98}, suggesting that studies of $\mu$ Jy radio sources are important for understanding the star formation history and evolution of ordinary galaxies \cite{haa00}."
 Whilst deep yoy field. surveys have made considerable progress in recent vears. no studies of radio sources in distant clusters have been carried out to comparable depth.," Whilst deep $\mu$ Jy field surveys have made considerable progress in recent years, no studies of radio sources in distant clusters have been carried out to comparable depth."
 At low redshift ες« 0.09). the radio luminosity function in clusters. when viewed in terms of the proportion of optical galaxies which have a given radio luminosity. is statistically incistinguishable from that of the field (2)..," At low redshift $z<0.09$ ), the radio luminosity function in clusters, when viewed in terms of the proportion of optical galaxies which have a given radio luminosity, is statistically indistinguishable from that of the field \cite{led96}."
 Phis is somewhat unexpected. since the onset of a starburst or AGN is likely to be induced. either by an infall of σας onto the galaxy or through a weak interaction or merger with a companion galaxy: it would be natural to assume that each of these phenomena would occur at a dillerent rate in cluster environments than in the field., This is somewhat unexpected since the onset of a starburst or AGN is likely to be induced either by an infall of gas onto the galaxy or through a weak interaction or merger with a companion galaxy: it would be natural to assume that each of these phenomena would occur at a different rate in cluster environments than in the field.
 In low redshift clusters the similarity may be due to the relatively relaxed states of the clusters. with much of the gas and galaxies Iving in stable virialised orbits and the gas having been stripped from the ealaxies in the cluster centre (e.g. Gunn Gott 1972): these ellects reduce interactions ancl prevent a higher fraction of radio sources forming in these environments.," In low redshift clusters the similarity may be due to the relatively relaxed states of the clusters, with much of the gas and galaxies lying in stable virialised orbits and the gas having been stripped from the galaxies in the cluster centre (e.g. Gunn Gott \nocite{gun72}: these effects reduce interactions and prevent a higher fraction of radio sources forming in these environments."
 Clusters. of ealaxies at higher redshifts. however. may. still be in their formation process. with relatively high galaxy merger rates (ο.g. van Dokkum and a plentiful supply of disturbed gas.," Clusters of galaxies at higher redshifts, however, may still be in their formation process, with relatively high galaxy merger rates (e.g. van Dokkum \nocite{dok99a} and a plentiful supply of disturbed gas."
 Ehese should provide ideal laboratories to induce starbursts ancl [GN (cf., These should provide ideal laboratories to induce starbursts and AGN (cf.
 the high fraction of post.starburst galaxies in z0.5 clusters: Poggianti et al Dwarakanath Owen (?). carried out a detailed racio study of two z0.25 ‘lusters. Abell 2125 ancl Abell 2645. which have similar redshifts ancl richnesses (Abell class 4). but very dilferent fractions of ButcherOenmler blue galaxies: Abell 2125 has a blue galaxy fraction of 0.19. whereas that of Abell 2645 is only 0.03.," the high fraction of post–starburst galaxies in $z \sim 0.5$ clusters; Poggianti et al \nocite{pog00}
 Dwarakanath Owen \shortcite{dwa99} carried out a detailed radio study of two $z \approx 0.25$ clusters, Abell 2125 and Abell 2645, which have similar redshifts and richnesses (Abell class 4), but very different fractions of Butcher–Oemler blue galaxies; Abell 2125 has a blue galaxy fraction of 0.19, whereas that of Abell 2645 is only 0.03."
 They found that the radio luminosity istribution of cluster members of Abell 2125 is bimocdal. with a peak at about 10757 1 composed. cntirels⋪⋅⇁ of AGN. and a second higher peak at or below (due to the detection limit) 10777 composed of a mixture of AGN and star-formüng galaxies (see Section 5.4 [or more details).," They found that the radio luminosity distribution of cluster members of Abell 2125 is bimodal, with a peak at about $10^{24.5}$ $^{-1}$ composed entirely of AGN, and a second higher peak at or below (due to the detection limit) $10^{22.5}$ $^{-1}$ composed of a mixture of AGN and star-forming galaxies (see Section \ref{lumdist} for more details)."
 En Abell 2645 many fewer radio sources were detected. with essentially the entire lower luminosity class of sources missing.," In Abell 2645 many fewer radio sources were detected, with essentially the entire lower luminosity class of sources missing."
 The presence of this lower luminosity class of racio sources therefore seems to be connected with 1e presence of the blue ButeherOemler galaxies. although 1© (wo. populations do not overlap strongly: only one of 10 blue galaxies in the Abell 2125 cluster has an associated radio source.," The presence of this lower luminosity class of radio sources therefore seems to be connected with the presence of the blue Butcher–Oemler galaxies, although the two populations do not overlap strongly: only one of the blue galaxies in the Abell 2125 cluster has an associated radio source."
 Smail ct al (72) have made a deep radio observation (rms L41072 7 in our assumed cosmology) of he z=0.41 cluster | 4713 and detect S cluster raciosources*.. associated. with a range of galaxy morphologies rom ellipticals through to Sdtype galaxies.," Smail et al \shortcite{sma99b} have made a deep radio observation (rms $1.4 \times 10^{22}$ $^{-1}$ in our assumed cosmology) of the $z=0.41$ cluster $+$ 4713 and detect 8 cluster radio, associated with a range of galaxy morphologies from ellipticals through to Sd–type galaxies."
 At higher redshifts. Stocke shortcitesto99b have studied radio galaxies in clusters with redshifts 0.3<2S0.8 to a limiting point source ο radio luminosity of 10757 5.," At higher redshifts, Stocke \\shortcite{sto99b} have studied radio galaxies in clusters with redshifts $0.3 < z \lta 0.8$ to a limiting point source GHz radio luminosity of $10^{23.5}$ $^{-1}$."
 Thev found no evidence for evolution of the population of radio sources between redshift 2~0.8 and the present epoch., They found no evidence for evolution of the population of radio sources between redshift $z \sim 0.8$ and the present epoch.
 However. their observations are only sensitive enough to detect the brighter of the two populations of cluster radio sources cliscovered by Dwarakanath Owen (2).. and it is the fainter population which appears to depend. upon the dynamical state of the cluster. and. hence may. be expected to show strong recdshilt evolution.," However, their observations are only sensitive enough to detect the brighter of the two populations of cluster radio sources discovered by Dwarakanath Owen \shortcite{dwa99}, and it is the fainter population which appears to depend upon the dynamical state of the cluster, and hence may be expected to show strong redshift evolution."
 The goal of the current project. therefore. is to investigate in detail the nature of the fainter radio source population in high redshift clusters.," The goal of the current project, therefore, is to investigate in detail the nature of the fainter radio source population in high redshift clusters."
 This aim was adcdressec by carrving out a deep VLA observation of the rich cluster 03 at z=O83., This aim was addressed by carrying out a deep VLA observation of the rich cluster $-$ 03 at $z=0.83$.
 The nature of this cluster anc the existing observations of it are described in Section 2.. and the new VLA observations are cescribed in Section 3..," The nature of this cluster and the existing observations of it are described in Section \ref{ms1054}, and the new VLA observations are described in Section \ref{vlaobs}. ."
 The radio source population is investigated and compare with the optical imaging cata in Section 4.., The radio source population is investigated and compared with the optical imaging data in Section \ref{radpop}.
 In Section 5 an analysis is mace of the cluster radio source population. including a detailed: comparison of the radio. optical anc emission line properties of the cluster galaxies.," In Section \ref{hostgals} an analysis is made of the cluster radio source population, including a detailed comparison of the radio, optical and emission line properties of the cluster galaxies."
 Conclusions, Conclusions
to infinity specifies a relation between the black hole mass AL and the parameters of the galaxy potential. particularly the velocity dispersion. ic. an AZ0 relation.,"to infinity specifies a relation between the black hole mass $M$ and the parameters of the galaxy potential, particularly the velocity dispersion, i.e. an $M - \sigma$ relation."
 For a general mass clistribution MiCI) we can use the first integral (20)) to do this., For a general mass distribution $M_{\rm tot}(R)$ we can use the first integral \ref{int}) ) to do this.
 However for a simple isothermal potential the equation of motion has the analytic solution where [4.605 are the position ancl speed. of the shell at time /=0 (ine. 2005).," However for a simple isothermal potential the equation of motion has the analytic solution where $R_0, v_0$ are the position and speed of the shell at time $t=0$ (King, 2005)."
 For large times the first term dominates. and the shell can reach arbitrarily large τας] if and only if the black hole mass exceeds the critical value This is very close to the observed. AM6 relation (4)) (cf Wing. 2005).," For large times the first term dominates, and the shell can reach arbitrarily large radii if and only if the black hole mass exceeds the critical value This is very close to the observed $M - \sigma$ relation \ref{msig}) ) (cf King, 2005)."
 At sulliciently large radii the quasar radiation Ποια is too cilute to cool the wind shock. and the shell accelerates bevond the escape value. cutting oll the galaxy. and. establishing the blackhole mass — bulgemass relation (cf Wine. 2003. 2005).," At sufficiently large radii the quasar radiation field is too dilute to cool the wind shock, and the shell accelerates beyond the escape value, cutting off the galaxy and establishing the black–hole mass – bulge--mass relation (cf King, 2003, 2005)."
" We see that the time for a continuouslydriven shell to reach a given radius 1072,50 kpe is For large Ao this time is significantly longer than the Salpeter time. iniplving that the black hole mass AJ must increase above the threshold value AZ, before the shell reaches large radii."," We see that the time for a continuously–driven shell to reach a given radius $R = 10R_{10}$ kpc is For large $R_{10}$ this time is significantly longer than the Salpeter time, implying that the black hole mass $M$ must increase above the threshold value $M_{\sigma}$ before the shell reaches large radii."
 Vhis may suggest that for galaxies with larec bulge raclii. the black hole mass mav tend to lie above the Alσ relation.," This may suggest that for galaxies with large bulge radii, the black hole mass may tend to lie above the $M- \sigma$ relation."
 There is some suggestion of this in the observational data (Marconi Llunt. 2003).," There is some suggestion of this in the observational data (Marconi Hunt, 2003)."
 Llowever the uncertainty here is in knowing just what radius the shell must reach in order to shut olf further accretion on to the black hole., However the uncertainty here is in knowing just what radius the shell must reach in order to shut off further accretion on to the black hole.
 We see from the reasoning of the last Section that the interaction beween the quasar wind and its host establishing heAJe relationis crucially ‘momentumciriven' rather han οποιονdriven., We see from the reasoning of the last Section that the interaction beween the quasar wind and its host establishing the $M - \sigma$ relation is – crucially – `momentum–driven' rather than `energy–driven'.
 This equivalent to requiring ellicient shock cooling., This equivalent to requiring efficient shock cooling.
 An energydriven shock (e.g. Silk Rees. 1998) would result in a much smaller. black bole mass for or a given σ than observed.," An energy–driven shock (e.g. Silk Rees, 1998) would result in a much smaller black hole mass for for a given $\sigma$ than observed."
 Instead. of the momentun rate Leaafe balancing the weight of swept.up gas σ.σ. which is what produces the momentum¢criven relation (27)). an energy.driven. shock would equate the energy deposition rate to the rate of working against this weight.," Instead of the momentum rate $\le/c$ balancing the weight of swept–up gas $4f_g\sigma^4/G$, which is what produces the momentum–driven relation \ref{msig2}) ), an energy–driven shock would equate the energy deposition rate to the rate of working against this weight."
 1n the near[Exddington regime the result is Le. which lies well below the observed. relation (4))., In the near–Eddington regime the result is i.e. which lies well below the observed relation \ref{msig}) ).
 The coupling adopted. in cosmological simulations evidently ensure that the interstellar medium feels. the outflow momentum rather than its energy. in addition to the “enerey ellicieney! ~4/22:0.05 noted above.," The coupling adopted in cosmological simulations evidently ensure that the interstellar medium feels the outflow momentum rather than its energy, in addition to the `energy efficiency' $\sim \eta/2 \simeq 0.05$ noted above."
 ] note finally that if instead. of the nearEddington regime considered here. an energy.αγνοία outflow had a large Eddington ratio H12271. we would have to replace (29)) by and thus (30)) by which is still smaller.," I note finally that if instead of the near–Eddington regime considered here, an energy–driven outflow had a large Eddington ratio $\dot m >>1$, we would have to replace \ref{work}) ) by and thus \ref{energysig}) ) by which is still smaller."
 Indeed one could imagine a situation in which the central black holes of medium or large galaxies obeved. (32)) rather than the observed. (4) and sellconsistently had central accretion rates well above I5ddington. since (6)) would now become The high optical depth implied bv the large Eddington ratio might) then prevent efficient Compton cooling. justifving the original hypothesis of energy.driven outflow.," Indeed one could imagine a situation in which the central black holes of medium or large galaxies obeyed \ref{energysig2}) ) rather than the observed \ref{msig}) ) and self–consistently had central accretion rates well above Eddington, since \ref{eddrat}) ) would now become The high optical depth implied by the large Eddington ratio might then prevent efficient Compton cooling, justifying the original hypothesis of energy–driven outflow."
 lt ds interesting that observation docs not sccm to eive examples of this possibilitv.. i.c. medium or large ealaxies with very lowmass central black holes which could accrete at. very high Exldington ratios.," It is interesting that observation does not seem to give examples of this possibility, i.e. medium or large galaxies with very low–mass central black holes which could accrete at very high Eddington ratios."
 Phe reason may be the inherent. tendency. of this lowmass sequence to. move irreversibly over time to the highmass case specified. by the usual Al6 relation (4)): thus steady sub. or nearEdedington accretion on to such a hole could. eventually. increase its mass to the point that it was unlikely to accrete at high Edeington ratios., The reason may be the inherent tendency of this low–mass sequence to move irreversibly over time to the high–mass case specified by the usual $M - \sigma$ relation \ref{msig}) ): thus steady sub– or near--Eddington accretion on to such a hole could eventually increase its mass to the point that it was unlikely to accrete at high Eddington ratios.
 This would then imply efficient cooling of the shock and thus à momentumcriven outflow., This would then imply efficient cooling of the shock and thus a momentum–driven outflow.
 On large scales the outflows described in Section 5.2 above all have (outer) shock velocities limited by the bulge velocity dispersion σ., On large scales the outflows described in Section 5.2 above all have (outer) shock velocities limited by the bulge velocity dispersion $\sigma$.
 Yet optical aud. UV. observations give clear evidence of outllows with velocities of several times this value., Yet optical and UV observations give clear evidence of outflows with velocities of several times this value.
 These cannot be the central quasar winds with ο~0.16 discussed in Section 3.," These cannot be the central quasar winds with $v
\sim 0.1c$ discussed in Section 3."
 Outllows confined. within a few parsees of the black hole may be the nearzone winds inside Riy discussed in Section 5.1. but high.velocity outllows are often seen or inferred on scales comparable with the entire ealaxy.," Outflows confined within a few parsecs of the black hole may be the near–zone winds inside $\rinf$ discussed in Section 5.1, but high–velocity outflows are often seen or inferred on scales comparable with the entire galaxy."
 Some of these outllows are seen in Compact radio sources (e.g. Llolt et al..," Some of these outflows are seen in compact radio sources (e.g. Holt et al.,"
 2008). some in Sevfert. galaxies which are clearly sub.Eddington. and others in post.starburst galaxies (cf Tremonti et al..," 2008), some in Seyfert galaxies which are clearly sub–Eddington, and others in post–starburst galaxies (cf Tremonti et al.,"
 2007)., 2007).
 The latter could result from the combined cllects of stellar winds and supernovae. but the known association of starbursts ancl AGN leave open the possibility that black holes may be the ultimate driver.," The latter could result from the combined effects of stellar winds and supernovae, but the known association of starbursts and AGN leave open the possibility that black holes may be the ultimate driver."
 There is a simple interpretation of such largescale high outllows., There is a simple interpretation of such large–scale high--velocity outflows.
" Consider a galaxy in which the SAIBIT has reached the mass Al, given by eqn (27)). with the"," Consider a galaxy in which the SMBH has reached the mass $M_{\sigma}$ given by eqn \ref{msig2}) ), with the"
the day and the night sides are equal.,the day and the night sides are equal.
 This is also a reasonable assumption for our purpose., This is also a reasonable assumption for our purpose.
 Budaj (2011)) calculated shapes and variations 1n the effective gravity over the surface of all currently known transiting exoplanets., Budaj \cite{budaj11}) ) calculated shapes and variations in the effective gravity over the surface of all currently known transiting exoplanets.
 He found that largest departures from the spherical symmetry are expected for WASP-12b and WASP-19b. of about12-15%.. respectively which translates into the variations of the gravity over the surface of about25-32%.," He found that largest departures from the spherical symmetry are expected for WASP-12b and WASP-19b, of about, respectively which translates into the variations of the gravity over the surface of about."
. The gravity difference between the substellar and antistellar points were less than2%., The gravity difference between the substellar and antistellar points were less than.
. For convenience. we assume that the thickness of the atmosphere is small compared to the planetary radius. so that the gravity 15 a constant.," For convenience, we assume that the thickness of the atmosphere is small compared to the planetary radius, so that the gravity is a constant."
 The entropy of the planetary core is thus determined by the entropy at the base of the convection zone. which in turn is determined by the magnitude of stellar irradiation. radiation transport in the atmosphere. and other details of the atmospheric physics.," The entropy of the planetary core is thus determined by the entropy at the base of the convection zone, which in turn is determined by the magnitude of stellar irradiation, radiation transport in the atmosphere, and other details of the atmospheric physics."
 Therefore. sophisticated evolutionary models have to take into account proper boundary conditions. in particular the values of the entropy and the radiatioαυ] cooling. using detailed atmospheric models. as was done e.g. in Burrows. Sudarsky. Hubbard (2003)). Burrows et al. (2007a)).," Therefore, sophisticated evolutionary models have to take into account proper boundary conditions, in particular the values of the entropy and the radiation cooling, using detailed atmospheric models, as was done e.g. in Burrows, Sudarsky, Hubbard \cite{bsh03}) ), Burrows et al. \cite{bhb07}) ),"
 Fortney. Marley. Barnes (2007)). and. Liu. Burrows. Ibgui (2008)).," Fortney, Marley, Barnes \cite{fmb07}) ), and Liu, Burrows, Ibgui \cite{lbi08}) )."
 The amount of radiation cooling from the day and the night sides is calculated in the following way: From the definition of the effective temperature. the total energy flux incoming at a unit surface area at the lower base of an atmosphere is given by F=oTiy.," The amount of radiation cooling from the day and the night sides is calculated in the following way: From the definition of the effective temperature, the total energy flux incoming at a unit surface area at the lower base of an atmosphere is given by $F= \sigma T_{\rm eff}^4$."
 Since no energy is being lost in the atmosphere. this quantity also represents the total radiation flux escaping from the unit surface at the upper boundary of the atmosphere.," Since no energy is being lost in the atmosphere, this quantity also represents the total radiation flux escaping from the unit surface at the upper boundary of the atmosphere."
 Therefore. it provides à quantitative measure of the radiator heat loss through the planet’s atmosphere. and. thus. from the whole planet.," Therefore, it provides a quantitative measure of the radiation heat loss through the planet's atmosphere, and, thus, from the whole planet."
 For usual atmospheric models. the effective temperature. together with surface gravity (and overall chemical composition) are taken as fundamental parameters. while from the point of view of interior and. evolutionary models. the primary parameter is the core entropy.," For usual atmospheric models, the effective temperature, together with surface gravity (and overall chemical composition) are taken as fundamental parameters, while from the point of view of interior and evolutionary models, the primary parameter is the core entropy."
 We divide the whole planetary surface into two parts. and treat the overall planetary atmosphere as composed of two distinct atmospheres — an averaged day-side atmosphere and an averaged night-side atmosphere.," We divide the whole planetary surface into two parts, and treat the overall planetary atmosphere as composed of two distinct atmospheres – an averaged day-side atmosphere and an averaged night-side atmosphere."
" The effective temperature on the day side is called 7. and on the night side 7,,."," The effective temperature on the day side is called $T_{d}$, and on the night side $T_{n}$."
" To avoid confusion with other possible meanings of the term ""effective temperature"" used for instance in planetary science. we will call Ty and 7,, “intrinsic effective temperatures”."," To avoid confusion with other possible meanings of the term “effective temperature” used for instance in planetary science, we will call $T_{d}$ and $T_{n}$ “intrinsic effective temperatures”."
 Obviously. these effective. temperatures should not be confused with brightness temperatures or an actual atmospheric temperature. which is a function of depth in the atmosphere and strongly depends on the irradiation.," Obviously, these effective temperatures should not be confused with brightness temperatures or an actual atmospheric temperature, which is a function of depth in the atmosphere and strongly depends on the irradiation."
 As mentioned above. to compute the total heat loss consistently. the model atmospheres for the day and the night side must have the same entropy at the convective base.," As mentioned above, to compute the total heat loss consistently, the model atmospheres for the day and the night side must have the same entropy at the convective base."
 To this end. we calculate a grid of models with/without the irradiation corresponding to the day/night side of the planet for a range of day- and night-side intrinsic effective temperatures and surface gravities (log e).," To this end, we calculate a grid of models with/without the irradiation corresponding to the day/night side of the planet for a range of day- and night-side intrinsic effective temperatures and surface gravities $\log g$ )."
 Each model has a certain entropy in the convection zone., Each model has a certain entropy in the convection zone.
" In the next step we match the entropy and gravity of the day and night sides. which results in different effective temperatures (7,#Τη) for the day and night sides."," In the next step we match the entropy and gravity of the day and night sides, which results in different effective temperatures $T_{d}\neq T_{n}$ ) for the day and night sides."
 Notice that we can obtain a unique match because entropy is a monotonic function of the intrinsic effective temperature., Notice that we can obtain a unique match because entropy is a monotonic function of the intrinsic effective temperature.
" Since T, and 7, represent the total radiation flux on the day and night sides. they represent the day and the night side internal heat loss. or. equivalently. the cooling of the interior."," Since $T_{d}$ and $T_{n}$ represent the total radiation flux on the day and night sides, they represent the day and the night side internal heat loss, or, equivalently, the cooling of the interior."
" The total internal heat loss (cooling). Less. from the planet is then given by where A, TAis the radius of the planet. ~ is the Stefan- constant. and Το is the composite intrinsic effective temperature."," The total internal heat loss (cooling), $L_{\rm cool}$ , from the planet is then given by where $R_{p}$ is the radius of the planet, $\sigma$ is the Stefan-Boltzmann constant, and $T_{\rm eff}$ is the composite intrinsic effective temperature."
 The individual model atmospheres are computed using the code designed to model atmospheres of irradiated giant planets and brown dwarfs., The individual model atmospheres are computed using the code designed to model atmospheres of irradiated giant planets and brown dwarfs.
 This code is a version of the stellar atmosphere code (Hubeny 1988:: Hubeny Lanz 1995)): modifications for are described in Hubeny. Burrows Sudarsky (2003)) and Sudarsky. Burrows Hubeny (2003)).," This code is a version of the stellar atmosphere code (Hubeny \cite{hubeny88}; Hubeny Lanz \cite{hl95}) ); modifications for are described in Hubeny, Burrows Sudarsky \cite{hbs03}) ) and Sudarsky, Burrows Hubeny \cite{sbh03}) )."
 The code solves the hydrostatic and the radiativetconvective equilibrium equations. assumes LTE. and can take into account clouds and departures from local chemical equilibrium (Hubeny Burrows 2007)).," The code solves the hydrostatic and the radiative+convective equilibrium equations, assumes LTE, and can take into account clouds and departures from local chemical equilibrium (Hubeny Burrows \cite{hb07}) )."
 The computed models represent separately averaged-day and averaged-night side atmospheres., The computed models represent separately averaged-day and averaged-night side atmospheres.
 The upper boundary ts set to the pressure 107 bars and extends deeply into the convectioi zone., The upper boundary is set to the pressure $10^{-5}$ bars and extends deeply into the convection zone.
 At a specified depth range. we allow for an energy removal on the day side. and an energy deposition on the night side of the planet. using the procedure described i1 Burrows. Buda) Hubeny (2008)).," At a specified depth range, we allow for an energy removal on the day side, and an energy deposition on the night side of the planet, using the procedure described in Burrows, Budaj Hubeny \cite{bbh08}) )."
" The amount of the heat redistribution between the day and night sides is parametrizec by the P, parameter. which is defined as the fraction of the incoming stellar irradiation that ts transferred from the day side to the night side and radiated from there (Burrows. Sudarsky Hubeny 2006))."," The amount of the heat redistribution between the day and night sides is parametrized by the $P_{n}$ parameter, which is defined as the fraction of the incoming stellar irradiation that is transferred from the day side to the night side and radiated from there (Burrows, Sudarsky Hubeny \cite{bsh06}) )."
 We consider here the well-known planet HD 209458b., We consider here the well-known planet HD 209458b.
" If not stated otherwise. we assume solar chemical composition of the planetary atmosphere. energy removal/deposition at 0.3 bar. and P,,=0.3."," If not stated otherwise, we assume solar chemical composition of the planetary atmosphere, energy removal/deposition at 0.03-0.3 bar, and $P_{n}=0.3$."
 TIO and VO opacity is not considered., TiO and VO opacity is not considered.
 Opacities are taken from Sharp Burrows (2007)). assuming chemical-equilibrium compositions with rainout but no cloud opacity.," Opacities are taken from Sharp Burrows \cite{sb07}) ), assuming chemical-equilibrium compositions with rainout but no cloud opacity."
 Kuruez (1993)) spectrum of the parent star HD 209458 is used as the source of irradiation., Kurucz \cite{kurucz93}) ) spectrum of the parent star HD 209458 is used as the source of irradiation.
 Parameters of the star and planet are taken from Henry et al. (2000)).," Parameters of the star and planet are taken from Henry et al. \cite{hmb00}) ),"
 Charbonneau et al. (20000).," Charbonneau et al. \cite{cbl00}) ),"
 and Knutson et al. (2007))., and Knutson et al. \cite{kcn07}) ).
 We assume Το=6000 K. Rz1.13RX. for the star. and a semi-major axis of the planet's orbit &=0.045 AU.," We assume $T_{\rm eff}=6000$ K, $R=1.13\ R_{\odot}$ for the star, and a semi-major axis of the planet's orbit $a=0.045$ AU."
 The models were calculated for a eric of effective temperatures ranging from 50 K up to 300 K anc gravities from 2.4 up to 3.6 (egs)., The models were calculated for a grid of effective temperatures ranging from 50 K up to 300 K and gravities from 2.4 up to 3.6 (cgs).
 We plot in Fig., We plot in Fig.
 | the entropy (per baryon. dividec by the Boltzmann constant k) as a function of temperature and pressure (Saumon. Chabrier. Van Horn 1995)).The crucial point is that the entropy increases with temperature anc decreases with pressure.," \ref{ent} the entropy (per baryon, divided by the Boltzmann constant $k$ ) as a function of temperature and pressure (Saumon, Chabrier, Van Horn \cite{sc95}) ).The crucial point is that the entropy increases with temperature and decreases with pressure."
 In Fig., In Fig.
 2. we plot the temperature- pressure (7-P) profiles as a function of the intrinsic effective, \ref{tefftp} we plot the temperature- pressure $T$ $P$ ) profiles as a function of the intrinsic effective
 eres (Farrell2009)..,$\sim$ $^{42}$ $^{-1}$.
" The recent discovery the optical IILX-1 speetim"" OF(W MD ierseiiaOnionet al.lino consistentin with JL, at the redshift aeE nr 2""[2.16m 0.0225)2l223) reve hbivvy0vocum ⋅⋅d ↕↓↸∖∩∏↑↴∖↴⊊∐⋅↑↴∖↴∩↕↑↕↓↸∖↸∖≺↕≦↴⋋↸∖⊣⋝∐∺∪⋜↧↴∖↴⋝∐⋅⋜↧↕≦↴⋋⋜↕↕⋜↧∑⋅↖⇁⊏⊱↻≥↓⋅≩≓D19 310 !. uNSIUSUSN DU Cu Meeorc ", The recent discovery in the optical HLX-1 spectrum of an emission line consistent with $H_{\alpha}$ at the redshift of ESO 243-49 $z=0.0223$ ) irrevocably confirms its association with this galaxy at a distance of 95 Mpc.
"! nWith observed MuX-rayu huninosities ""reaching above 10 1 |TILN-1 is super-Eddingtow if the black-hole’s niv is less than ~ 10!AL.."," With observed X-ray luminosities reaching above $10^{42}$ $^{-1}$ HLX-1 is super-Eddington if the black-hole's mass is less than $\sim10^{4}\,\Msun$."
.Beauing effects (e.g. −∖∖↕↾⋅⊓∣∣⊓∣∣⋯∙↧∣∪∣⋝∖≺⋅↾⋅∖⊒⋯⊳↾⋅∙∖⋅⊳∣⊒↕↜∟≯∙∣≺∙∣∣⊓∣∣∣⊒∐∣⇂∣∣∣∖≺∙↾⋅⋟∖∣↕∙∖⋅∎⇂∣∡∪↾⋅∣⊒↕ ∙ i ↴ ∙−↓∎∣≯∐↓↾∖⋮↕∩∣∣ AR, Beaming effects have been proposed as viable mechanisms for producing the apparent super-Eddington luminosities seen from other ULXs.
"ON?sepIi. dona ot7 oe dosTMoses. *CIE ΗΕ ΕΕ enebnman remeosseran e ΙΑπο ο esetμμ ologieu vite BAesew fro pes Omομμpicsecedroladiues dslack hole Departinentof aite et?Hes Hote Ul mdπό uthe jet-axis)and the nuiuositv er the IZ, line"," However, beaming is unlikely to explain HLX-1's extreme luminosity due to the observed large-scale variability (which appears similar to that seen from Galactic stellar mass black hole binaries that are not viewed down the jet-axis) and the luminosity of the $H_{\alpha}$ line"
"the limiting case of β—0, the Alfvénnic Mach number is given by Ma,=5)/[2(4—X)]VX(X+ for a perpendicular shock and Maj=VX for a parallel shock (e.g., Vrsnnak et al.","the limiting case of $\beta\rightarrow 0$, the Alfvénnic Mach number is given by $M_{{\rm A}\perp}=\sqrt{X(X+5)/[2(4-X)]}$ for a perpendicular shock and $M_{{\rm A}\parallel}=\sqrt{X}$ for a parallel shock (e.g., Vršnnak et al."
 2002)., 2002).
" The angles 0p, have been estimated pixel by pixel along the shock surfaces from the LASCO images by assuming a radial magnetic field above ~2Re..", The angles $\theta_{Bn}$ have been estimated pixel by pixel along the shock surfaces from the LASCO images by assuming a radial magnetic field above $\sim 2$.
" By taking, as a first order approximation for the more general case of an oblique shock, Ma,=(Ma.sin@g,)*+(Majcosgn), implying Ma,2May, we finally obtained the values plotted in Figure 2 "," By taking, as a first order approximation for the more general case of an oblique shock, $M_{{\rm A}\angle}=\sqrt{(M_{{\rm A}\perp}\sin\theta_{Bn})^2+(M_{{\rm A}\parallel}\cos\theta_{Bn})^2}$, implying $M_{{\rm A}\perp} \geq M_{{\rm A}\angle} \geq M_{{\rm A}\parallel}$, we finally obtained the values plotted in Figure 2 )."
"panels). Ma, is seen to maximizes at the shock center at both times, reaching a maximum value of ΜΑ,~1.8 at ~2.6Ro."," $M_{{\rm A}\angle}$ is seen to maximizes at the shock center at both times, reaching a maximum value of $M_{{\rm A}\angle}\simeq 1.8$ at $\sim 2.6$."
". Critical values for the Alfvénnic Mach number MX for collisionless non-relativistic shocks are provided by Edminston Kennel (1984), recently reviewed by Treumann (2009)."," Critical values for the Alfvénnic Mach number $M_{{\rm A}}^\star$ for collisionless non-relativistic shocks are provided by Edminston Kennel (1984), recently reviewed by Treumann (2009)."
" The above authors provide a theoretical MX curve for a fast-mode shock as a function of the angle 6p, for 8=0.", The above authors provide a theoretical $M_{{\rm A}}^\star$ curve for a fast-mode shock as a function of the angle $\theta_{Bn}$ for $\beta = 0$.
" From the 63, values derived along the shock front, we estimated Mx at different latitudes."," From the $\theta_{Bn}$ values derived along the shock front, we estimated $M_{{\rm A}}^\star$ at different latitudes."
" The resulting Mx variations along the front are seen to be opposite to those of ΜΑ,.", The resulting $M_{{\rm A}}^\star$ variations along the front are seen to be opposite to those of $M_{{\rm A}\angle}$.
" That is, Mx is seen to minimize at the shock center and maximize at the flanks."," That is, $M_{{\rm A}}^\star$ is seen to minimize at the shock center and maximize at the flanks."
" Comparison between the M4, and Mx curves in Figure 2 shows that, at earlier times and lower altitudes, Ma,>Mx at the center of the shock and Ma,«€Mx at the shock flanks, while, at later times and higher altitudes, Ma,<Mx at all latitudes."," Comparison between the $M_{{\rm A}\angle}$ and $M_{{\rm A}}^\star$ curves in Figure 2 shows that, at earlier times and lower altitudes, $M_{{\rm A}\angle} > M_{{\rm A}}^\star$ at the center of the shock and $M_{{\rm A}\angle} < M_{{\rm A}}^\star$ at the shock flanks, while, at later times and higher altitudes, $M_{{\rm A}\angle} < M_{{\rm A}}^\star$ at all latitudes."
" Hence, at earlier times the shock center is super-critical, making it a probable source for SEP acceleration."," Hence, at earlier times the shock center is super-critical, making it a probable source for SEP acceleration."
state spectra are shown in Fig.,state spectra are shown in Fig.
 2. with a binning of 50mA.," \ref{spectra_01,06} with a binning of 50."
 It can be seen that the 2006 spectrum shows much better resolved. absorption (roughs and was used in order to obtain the warm absorber parameters., It can be seen that the 2006 spectrum shows much better resolved absorption troughs and was used in order to obtain the warm absorber parameters.
 The fitting procedure follows our ion-bv-ion fitting method (Beharetal.2001:SakoBeharetal.2005).," The fitting procedure follows our ion-by-ion fitting method \citep{behar01, sako01, behar03, tomer05}."
. First. we fit for the broad-band continuumn.," First, we fit for the broad-band continuum."
 Subsequently. we fit the absorption features using template ionic spectra that include all of the absorption lines and photoelectric edges of each ion. but vary. with the broadening (so-called turbulent) velocity and the ionic column density.," Subsequently, we fit the absorption features using template ionic spectra that include all of the absorption lines and photoelectric edges of each ion, but vary with the broadening (so-called turbulent) velocity and the ionic column density."
 Covering factor of unity is used throughout this process., Covering factor of unity is used throughout this process.
 The “black” troughis of the leading lines of Ὁ and * strongly support this assumption.," The ""black"" troughs of the leading lines of $^{+6}$ and $^{+7}$ strongly support this assumption."
" Strong, emission lines are fitted as well.", Strong emission lines are fitted as well.
 The emission lines ave added alter the absorption components were modelled (hence. the emission lines are not absorbed).," The emission lines are added after the absorption components were modelled (hence, the emission lines are not absorbed)."
 Another interesting feature that can be seen in Fig 2. is the softening of the spectrum at higher flux levels.," Another interesting feature that can be seen in Fig \ref{spectra_01,06} is the softening of the spectrum at higher flux levels."
 This was already observed for NGC 3183 (Netzeretal.2003). and in fact expected [rom the cooling of a comptonizing corona above the accretion disc (ILaadtetal.2001)., This was already observed for NGC 3783 \citep{netzer03} and in fact expected from the cooling of a comptonizing corona above the accretion disc \citep{haardt01}.
. The continuum sspectrum of most ACNs can be characterized by a high-energv power-law aud a soft excess that rises above the power-law below ~1 keV. This soft excess is offen modeled with a blackbody. or modified blackbodsy. although it clearly is more spectrally complex aud possibly includes atomic features.," The continuum spectrum of most s can be characterized by a high-energy power-law and a soft excess that rises above the power-law below $\sim$ 1 keV. This soft excess is often modeled with a blackbody, or modified blackbody, although it clearly is more spectrally complex and possibly includes atomic features."
 For the 2006 spectra. we used a power law with a photon spectral index of D = 1.48. which was fitted to the 26 bband. aud a blackbody temperature of kL = 110 eV. This is a rather flat slope. which could be due to the band in which it is fitted. but it still provides an good fit to the spectrum (see Fig. 4)).," For the 2006 spectra, we used a power law with a photon spectral index of $\Gamma$ = 1.48, which was fitted to the 2–6 band, and a blackbody temperature of $kT$ = 110 eV. This is a rather flat slope, which could be due to the band in which it is fitted, but it still provides an good fit to the spectrum (see Fig. \ref{ngcfit}) )."
 The intensity spectrum Z;;(p) around an atomic absorption line /—j can be expressed as: where /j(77) represents the unabsorbed continuum intensity. σι”) denotes the line absorption," The intensity spectrum $I_{ij}(\nu )$ around an atomic absorption line $i
\rightarrow j$ can be expressed as: where $I_0(\nu)$ represents the unabsorbed continuum intensity, $\sigma_{ij}(\nu)$ denotes the line absorption"
diffuse X-rays around Terzan 5 which exteneded up 10 pe (Eger. Domainko Clapson 2010).,"diffuse X-rays around Terzan 5 which exteneded up $\sim10$ pc (Eger, Domainko Clapson 2010)."
 Although a clear scenario cannot be identified vet. these diffuse N-ravs are more likely io have a non-thermal origin (Eger. Domainko Clapson 2010).," Although a clear scenario cannot be identified yet, these diffuse X-rays are more likely to have a non-thermal origin (Eger, Domainko Clapson 2010)."
 Assuming these N-ravs are originated from the tail of ICs. the corresponding 5—ray spectrum can be caleulated (Cheng el al.," Assuming these X-rays are originated from the tail of ICS, the corresponding $\gamma-$ ray spectrum can be calculated (Cheng et al."
 2010)., 2010).
 Therefore. a svstematic search for the extended X-ray. and radio feature outside the hall-mass radii of the other 5--ταν GC's can provide us indpendent constraints.," Therefore, a systematic search for the extended X-ray and radio feature outside the half-mass radii of the other $\gamma-$ ray GCs can provide us indpendent constraints."
 In exploring the fundamental plane relations. our analvsis suggests that by combining the soft photon energy. densities with DP./|Fe/1l] the data scattering can be reduced.," In exploring the fundamental plane relations, our analysis suggests that by combining the soft photon energy densities with $\Gamma_{c}$ $\left[{\rm Fe/H}\right]$ the data scattering can be reduced."
 These best-fit relations can provide the indicators in identilving what kind of GCs are potential —yraw sources for a further search., These best-fit relations can provide the indicators in identifying what kind of GCs are potential $\gamma-$ ray sources for a further search.
 And the other wax round. any deeper 5—rav search [rom the GCs can result in an enlarged sample size ancl a lower sensitivity limit than (he current value (ie. 6x10.P erg ? +). which will certainly enable a further test for all these reported relations.," And the other way round, any deeper $\gamma-$ ray search from the GCs can result in an enlarged sample size and a lower sensitivity limit than the current value (i.e. $6\times10^{-12}$ erg $^{-2}$ $^{-1}$ ), which will certainly enable a further test for all these reported relations."
predict how the elliciency would depend on the spectrum.,predict how the efficiency would depend on the spectrum.
 In this context. a more detailed comparison of heat transfer in initially laminar versus initiallv turbulent svstems would be ol interest.," In this context, a more detailed comparison of heat transfer in initially laminar versus initially turbulent systems would be of interest."
 Financial support for this project was provided by the Space Telescope Science Institute erants HIST-ABR-11251.01-À and HIST-AB-12128.01-À: by the National Science Foundation under award AST-O807363 and NSF PIIY0903797: by the Department of Energy. under award DI-SCOO001063: and by Cornell University. grant. 41843-7012., Financial support for this project was provided by the Space Telescope Science Institute grants HST-AR-11251.01-A and HST-AR-12128.01-A; by the National Science Foundation under award AST-0807363 and NSF PHY0903797; by the Department of Energy under award DE-SC0001063; and by Cornell University grant 41843-7012.
 We wish to thank Jonathan Carroll. Iris Yirak ancl Brandon Shrover for useful discussions.," We wish to thank Jonathan Carroll, Kris Yirak and Brandon Shroyer for useful discussions."
 The MIID solver and the linear thermal diffision solver are verified by well-known tests such as the field loop convection problem and (he Guassian diffusion problem separately., The MHD solver and the linear thermal diffusion solver are verified by well-known tests such as the field loop convection problem and the Guassian diffusion problem separately.
 As a comprehensive test that involves both ATID and thermal diffusion. we use the magneto-thermal instability (AIT) problem to test the accuracy of the ASTRODEAR code with anisotropic heat conduction (Parrish&Stone (2005)..Cunninghametal. (2009))).," As a comprehensive test that involves both MHD and thermal diffusion, we use the magneto-thermal instability (MTI) problem to test the accuracy of the ASTROBEAR code with anisotropic heat conduction \citet{par05}, \citet{cun09}) )."
 The problem involves setting up a 2-D temperature profile wilh uniform gravity pointing on the v direction., The problem involves setting up a 2-D temperature profile with uniform gravity pointing on the y direction.
 The domain is square with length of 0.1 in compuational units., The domain is square with length of $0.1$ in compuational units.
 The temperature and density. profiles are:, The temperature and density profiles are:
Ilstoricallw. the evolution of the research in quantum eravity can roughly be divided in five periods.,"Historically, the evolution of the research in quantum gravity can roughly be divided in five periods."
discussed by ο who use a theoretical stellar masseracdius relation to break this degeneracy.,discussed by \citet{sea03a} who use a theoretical stellar mass-radius relation to break this degeneracy.
 We provide a simplified treatment here to illustrate the nature of the degeneraey and how it can be broken with observations of transit timing variations., We provide a simplified treatment here to illustrate the nature of the degeneracy and how it can be broken with observations of transit timing variations.
 Consider a planetary system with two transiting planets on circular orbits which are coplanar. exactly edge-on. and have measured radial velocity amplitudes.," Consider a planetary system with two transiting planets on circular orbits which are coplanar, exactly edge-on, and have measured radial velocity amplitudes."
 Well assume that the star is uniform in surface brightness and that mij.m»«mo.," We'll assume that the star is uniform in surface brightness and that $m_1, m_2 \ll m_0$."
 We also assume that the unperturbecl periods £1.75 can be measured from the duration between transits.," We'll also assume that the unperturbed periods $P_1, P_2$ can be measured from the duration between transits."
" Phen there are eight ohvsical parameters of interest which describe the svstem: mo.mg.Ry.Iso,. and es where 5; ave the radii of the star and planets."," Then there are eight physical parameters of interest which describe the system: $m_0, m_1, m_2, R_0, R_1, R_2, a_1,$ and $a_2$ where $R_i$ are the radii of the star and planets."
 Without measuring the transit timing variations. there are a total of ten parameters which can be measured: Wy.νοepppolatyeAL.AbDES. and {νι where fr; labels the duration of transit [rom mid-ingress to mid-egress. {η abels the curation of ingress or egress for planet j. Av; are the velocity amplitudes of the two planets. and AL) are the relative depths of the transits in units of the uneclipsed brightness of the star (for planet j).," Without measuring the transit timing variations, there are a total of ten parameters which can be measured: $K_1, K_2, t_{T1}, t_{T2}, t_{g1}, t_{g2}, \Delta F_1,
\Delta F_2, P_1,$ and $P_2$, where $t_{Tj}$ labels the duration of transit from mid-ingress to mid-egress, $t_{gj}$ labels the duration of ingress or egress for planet $j$, $K_j$ are the velocity amplitudes of the two planets, and $\Delta F_j$ are the relative depths of the transits in units of the uneclipsed brightness of the star (for planet $j$ )."
 Although there are more constraining xwameters than model parameters. there is a degeneracy since some of the observables are redundant.," Although there are more constraining parameters than model parameters, there is a degeneracy since some of the observables are redundant."
 All of the svsten xwameters can be expressed in terms of observables and the ratio of the mass to radius of the star. mofi. where j=1.2 labels cach planet.," All of the system parameters can be expressed in terms of observables and the ratio of the mass to radius of the star, $m_0/R_0$, where $j=1,2$ labels each planet."
 From this information alone one can constrain the density of the star (2).., From this information alone one can constrain the density of the star \citep{sea03a}.
 For the simplified case discussed here. [or either planet. (thisdillerssligthlv[romtheexpressionin7.sincewedefinetransitdurationfrommid in/egress)..," For the simplified case discussed here, for either planet \citep[this differs sligthly from the expression in][since we define
the transit duration from mid in/egress]{sea03a}. ."
 Lf in addition. one can measure the amplitude of the transit timing variations of the outer planet. σοι then this determines the mass ratio.," If, in addition, one can measure the amplitude of the transit timing variations of the outer planet, $\sigma_2$, then this determines the mass ratio."
 For the case that the stars motion dominates the transit timing. For other cases. the transit timing amplitude can be computed numerically.," For the case that the star's motion dominates the transit timing, For other cases, the transit timing amplitude can be computed numerically."
 Phen. from the above expression for mi;fry one can find the ratio of the mass to the radius of the star Combined with the measurement of the density. this gives the absolute mass and radius of the star.," Then, from the above expression for $m_i/m_0$ one can find the ratio of the mass to the radius of the star Combined with the measurement of the density, this gives the absolute mass and radius of the star."
 This procedure requires no assumptions about the mass-racius relation for the host star. and in principle could be used to measure this relation.," This procedure requires no assumptions about the mass-radius relation for the host star, and in principle could be used to measure this relation."
 Hf one can also measure transit timing variations for the inner planet. then an extra constraint can be obtained where f(a) is a function derived from averaging equation AT... (," If one can also measure transit timing variations for the inner planet, then an extra constraint can be obtained where $f(\alpha)$ is a function derived from averaging equation \ref{sigcirc1}. ("
Note that the phase of the orbits is needed for this equation. whieh can be found from the velocity amplitude curve).,"Note that the phase of the orbits is needed for this equation, which can be found from the velocity amplitude curve)."
 This provides an extra constraint on the system. and thus will be a check that this procedure is robust.," This provides an extra constraint on the system, and thus will be a check that this procedure is robust."
 Clearly we have made some drastically simplifving assumptions which are not true for any physical transit., Clearly we have made some drastically simplifying assumptions which are not true for any physical transit.
 Phe inclination of the orbits must be solved for. which can be done from the ratio of the durations of the ingress and transit and the change in ες. as demonstrated by ?..," The inclination of the orbits must be solved for, which can be done from the ratio of the durations of the ingress and transit and the change in flux, as demonstrated by \citet{sea03a}."
 In addition. limb-darkening must be included. and can be solved for with high signal-to-noise cleda as demonstrated by 2..," In addition, limb-darkening must be included, and can be solved for with high signal-to-noise data as demonstrated by \citet{bro01}."
 Finally. the orbits are not generally circular. so the parameters ej.ze;Q;.0m;. which can be derised [rom the velocity. amplitude measurements. should. be accounted for.," Finally, the orbits are not generally circular, so the parameters $e_j, \varpi_j, \Omega_j, \sigma_j$, which can be derived from the velocity amplitude measurements, should be accounted for."
 The general solution is rather complicated and. wotuel best be accomplished numerically. but the cegeneracy has a similar nature to the circular case and canin principle be broken bv the transit timing variations.," The general solution is rather complicated and would best be accomplished numerically, but the degeneracy has a similar nature to the circular case and canin principle be broken by the transit timing variations."
in Fig.,in Fig.
 3. degrade o those in Fig., \ref{fig:RV_wav_120000} degrade to those in Fig.
 5 (R=120.000).," \ref{fig:RV_wav_cal_120000} $\rm{R}=120,000$ )."
 Two sceratios of calibration are cousidered: Superimiposiug (Dotted). aud Nou-Common Path auc Bracketing (Solid).," Two scenarios of calibration are considered: Superimposing (Dotted), and Non-Common Path and Bracketing (Solid)."
 The dillerence between these two is whetler the S/N depends on he stellar flux., The difference between these two is whether the S/N depends on the stellar flux.
 In the 5iperiniposiug case. because the absorption cell is in the light yath of tlje stellar flux. the continuu of the resulting Locline absorption spectruu is determined by the he contiuuunm flux of a star.," In the Superimposing case, because the absorption cell is in the light path of the stellar flux, the continuum of the resulting Iodine absorption spectrum is determined by the the continuum flux of a star."
 Consequently. the RV calibration uncertaiuly is strongly cepeudent ou the incoming stellar ix.," Consequently, the RV calibration uncertainty is strongly dependent on the incoming stellar flux."
 Iu the comiparisou of the two Cases in Fig. 5..," In the comparison of the two cases in Fig. \ref{fig:RV_wav_cal_120000},"
 we see the Non-Commo1 Path and the Bracketi& methods always introduce less uucerainty in IRV calibration than the Stperimposiug method., we see the Non-Common Path and the Bracketing methods always introduce less uncertainty in RV calibration than the Superimposing method.
 TIe major reasou lor that is the S/N iu the former case may be optimized by adjusting the source intesity (Noun-Common Path) or the exposure ime (Bracketing)., The major reason for that is the S/N in the former case may be optimized by adjusting the source intensity (Non-Common Path) or the exposure time (Bracketing).
 The main e«oiclusion. about optimal observational bandpass frou. 83.2.1 renialus uncliaugec., The main conclusion about optimal observational bandpass from \ref{sec:OptiBand_Star} remains unchanged.
 The optimal baud for Doppler measureuments is in the NIR (Ix baud in paricular) for stars with specral types later than M65 from previous discussions iu this paper., The optimal band for Doppler measurements is in the NIR (K band in particular) for stars with spectral types later than M5 from previous discussions in this paper.
 However. oue iniportant element is nulssing in the diseussiou. which is the contamination [rou the telluric lines in the Earth's amosphere. which is a severe problem iu the NIR observation.," However, one important element is missing in the discussion, which is the contamination from the telluric lines in the Earth's atmosphere, which is a severe problem in the NIR observation."
 The quaritative analysis oL telluric line contamination is iutroduced in 82.5 aud we apply that inehod iu estimating the RV unceraiuty brought by the telluric contamination., The quantitative analysis of telluric line contamination is introduced in \ref{sec:Telluci} and we apply that method in estimating the RV uncertainty brought by the telluric contamination.
 We coufine our discussions for late-type M dwarls since NIR observation does not gain advantage for other spectral types earler {man M»V. Fig., We confine our discussions for late-type M dwarfs since NIR observation does not gain advantage for other spectral types earlier than M5V. Fig.
 6 shows au examdle of how RV uncertainty for au 19V star changes wih observational oiudpass uucder different vaues ofa (Le.. level of telluric liie removal. see Equation CL).," \ref{fig:RV_wav_cal_atm_120000} shows an example of how RV uncertainty for an M9V star changes with observational bandpass under different values of $\alpha$ (i.e., level of telluric line removal, see Equation \ref{eq:B_atm}) ))."
 1 inelicates 10 telluric liue remOVal alic Q indicates complete removal of elluric lines (see Ecuation (1)))., 1 indicates no telluric line removal and 0 indicates complete removal of telluric lines (see Equation \ref{eq:B_atm}) )).
 RV luctuation of LO niesC due to random aliospheric movement is assuiued i the calclation., RV fluctuation of 10 $\rm{m}\cdot\rm{s}^{-1}$ due to random atmospheric movement is assumed in the calculation.
 Bracketing RV calibra is assumed in the calculation., Bracketing RV calibration is assumed in the calculation.
 There are several points worth noting u this plot: 1. different observatioua bandpasses are alleced dillereutly by elluric lines. the significauce of tellurie line contamination is indicated by the span of RV uucertaities at different a values.," There are several points worth noting in this plot: 1, different observational bandpasses are affected differently by telluric lines, the significance of telluric line contamination is indicated by the span of RV uncertainties at different $\alpha$ values."
 For example. B band is the leas allectecd by tellurie lines because the RV ncertaiuties in B yan at different levels of telluric line remov:il remain roughly tje same. While J. A and A baucls suffer severe telluri: line contamination becatine any sinall chanee of a results in siguificaut change of RV uncertainty.," For example, $B$ band is the least affected by telluric lines because the RV uncertainties in $B$ band at different levels of telluric line removal remain roughly the same, while $J$, $H$ and $K$ bands suffer severe telluric line contamination because any small change of $\alpha$ results in significant change of RV uncertainty."
 2. If there is no attempt of removing telluric lijes [roin observed stellar spectrum (purple in Fig. 6)).," 2, If there is no attempt of removing telluric lines from observed stellar spectrum (purple in Fig. \ref{fig:RV_wav_cal_atm_120000}) ),"
 there is uo advantage in c)bserving late-type A dwarfs in NIR. RV uncertaiuty Is ominated by Earls atmosphere behavior in the NIB.," there is no advantage in observing late-type M dwarfs in NIR, RV uncertainty is dominated by Earth's atmosphere behavior in the NIR."
 Lu this case. the optimal baud is V. aud R baud.," In this case, the optimal band is $V$ and $R$ band."
 Only when a 0.01. ie. more than telluric line sreneth is removed. the advantage oL observing late-type M dwarls in the NIB. becomes obvious. al a factor of 73) improvement.," Only when $\alpha\leq$ 0.01, i.e., more than telluric line strength is removed, the advantage of observing late-type M dwarfs in the NIR becomes obvious, at a factor of $\sim$ 3 improvement."
 Iu practice. there lave been several examples in which telluric liue mocleling aud removal is demioustrated to be successful.," In practice, there have been several examples in which telluric line modeling and removal is demonstrated to be successful."
 ? achieved maximum ¢eviatious of less than and RMS deviations of less than with 42-2000 aud S/N>100 using a telluric standard star nearby, \citet{Vacca2003} achieved maximum deviations of less than and RMS deviations of less than with $R$ =2000 and $\geq$ 100 using a telluric standard star nearby
 005001180to them η was higher.,layer (as judged by how close the index of refraction at 500nm is to the bulk index) was consistently higher.
" Furthermore, on Wayelengttipτα) delta-doped devices, sputtering caused a variety of veproblems."," Furthermore, on tests with live delta-doped devices, sputtering caused a variety of problems."
 We believe interactions of the film material with the surface Si , We believe interactions of the film material with the surface Si created highly absorptive silicates.
qur later Os. We O@cssary η haz," As such, our later work has used ALD exclusively for $_2$ and $_2$ $_3$."
obeen nr bnce the préeursors , We plan to transition to ALD for MgO and $_2$ once the necessary precursors have been obtained.
"A forMicoming paper will discuMWavglengtlgmm) detail the comparisons between films made by ALD and sputtering (Greer, 2011 (submitted))."," A forthcoming paper will discuss in more detail the comparisons between films made by ALD and sputtering (Greer, 2011 (submitted))."
" AR coating films have shown good stability thus far and long term stability tests are underway, to be reported separately."," AR coating films have shown good stability thus far and long term stability tests are underway, to be reported separately."
 Tests remain to determine the feasibility of applying different non-overlapping layers to the same CCD., Tests remain to determine the feasibility of applying different non-overlapping layers to the same CCD.
 Figure 13 depicts the overall outlook for multiple coatings on a single CCD., Figure \ref{fig:ave} depicts the overall outlook for multiple coatings on a single CCD.
 Each film is shown over the wavelength range of interest., Each film is shown over the wavelength range of interest.
 The solid bar is an average of the reflectance values over that range., The solid bar is an average of the reflectance values over that range.
 This provides an encouraging outlook for a future detector., This provides an encouraging outlook for a future detector.
 Our team has applied the films discussed above to thinned and delta-doped CCDs., Our team has applied the films discussed above to thinned and delta-doped CCDs.
" Work on this is ongoing and another publication describes our results in detail, but it merits a short mention here."," Work on this is ongoing and another publication describes our results in detail, but it merits a short mention here."
" The procedure for thinning anddelta-doping, as described in Hoenk, (?),, was carried out in the MicroDevices Lab at JPL on both Cassini CCDs and Electron Multiplying CCDs (EMCCDs) (?).."," The procedure for thinning anddelta-doping, as described in Hoenk, \citep{2009Hoenk}, was carried out in the MicroDevices Lab at JPL on both Cassini CCDs and Electron Multiplying CCDs (EMCCDs) \citep{2001Mackay}. ."
 The, The
sources is such that unidentified sources make up a sienificaut portion of the EGRD.,sources is such that unidentified sources make up a significant portion of the EGRB.
" We use observatious of the ECGRD to deteriiue by what value of F5, the extrapolated fux distribution power-lawust break. so that the background is not overpredicted at auv energy range."," We use observations of the EGRB to determine by what value of $F_{\rm 8,min}$ the extrapolated flux distribution power-law break, so that the background is not overpredicted at any energy range."
 To eusure that we are placing the most conservative constraint tthat we allow extrapolation to the lowest fluxes possible). we will use the most generous observational deteriuimation of the ECRB. which is that of Sreekiunar et ((1998).," To ensure that we are placing the most conservative constraint that we allow extrapolation to the lowest fluxes possible), we will use the most generous observational determination of the EGRB, which is that of Sreekumar et (1998)."
 The requireineut that this background is not exceeded for any euerev above LOO MeV. returns the limiting breaking Huxes indicated by the three vertical lines in reflosuloss.., The requirement that this background is not exceeded for any energy above 100 MeV returns the limiting breaking fluxes indicated by the three vertical lines in \\ref{lognlogs}.
 Our results then indicate that. for auv of our three samples. we would only need to extrapola5 he cumulative fux distribution for at most an order of magnitude below the lower linut of the resolved fiux range to have unresolved unideutified sources comprise uost of the EGRD. at least at low energies.," Our results then indicate that, for any of our three samples, we would only need to extrapolate the cumulative flux distribution for at most an order of magnitude below the lower limit of the resolved flux range to have unresolved unidentified sources comprise most of the EGRB, at least at low energies."
 This is rot an extreme extrapolation. aud therefore a siguificaut contribution bv the “wuideutified™” class to the diffuse vackeround is likely.," This is not an extreme extrapolation, and therefore a significant contribution by the “unidentified” class to the diffuse background is likely."
 Iu order to interpret this result correctly. we should assess the qualitative and quantitative uncertainties associated with it.," In order to interpret this result correctly, we should assess the qualitative and quantitative uncertainties associated with it."
 First of all. a power-law extrapolation of the flux distribution to lower. unresolved flux values even for oulv one order of magnitude in flux is by no means seltevidenut. although it is the simplest and most straight-forward asstuuption.," First of all, a power-law extrapolation of the flux distribution to lower, unresolved flux values even for only one order of magnitude in flux is by no means self-evident, although it is the simplest and most straight-forward assumption."
 For example. the Denier (2007) fits to the flux distribution of resolved blazars (using a imodel based on the cosguological evolution of black hole jet sources) increase less steeply than a power law even very close to the EGRET scusitivity limit.," For example, the Dermer (2007) fits to the flux distribution of resolved blazars (using a model based on the cosmological evolution of black hole jet sources) increase less steeply than a power law even very close to the EGRET sensitivity limit."
 It is conceivable that unresolved members of the extragalactic unidentified source class could have a comparable cosmic evolution. aud a simular behavior iu the flux fiction.," It is conceivable that unresolved members of the extragalactic unidentified source class could have a comparable cosmic evolution, and a similar behavior in the flux function."
 Iu fact. our tight coustraiuts on the power-law extrapolation of the flux fiction may be taken to imply just that: the cosmological evolution of these sources to be such that the dux function deviates from a power-law form very fast. to eusure than unresolved members of this class do not overproduce the EGRB.," In fact, our tight constraints on the power-law extrapolation of the flux function may be taken to imply just that: the cosmological evolution of these sources to be such that the flux function deviates from a power-law form very fast, to ensure than unresolved members of this class do not overproduce the EGRB."
 Second. statistical aud svstematic unnucertaiuties iu the power-law fit» to the cuuulative flux function lave to be taken iuto account.," Second, statistical and systematic uncertainties in the power-law fits to the cumulative flux function have to be taken into account."
 In our case. the systematic uncertainties associated with the sample selection aud the small civnamical range in flux over which the flux function is sampled cominate over the statistical uncertainties in the power-law slope derived iu the selected fiux rauge. shown in Table 1..," In our case, the systematic uncertainties associated with the sample selection and the small dynamical range in flux over which the flux function is sampled dominate over the statistical uncertainties in the power-law slope derived in the selected flux range, shown in Table \ref{sometable}."
 The most extreme case is the oue where wuidentified sources are considered as a unified sample with identified blazars outside the mass., The most extreme case is the one where unidentified sources are considered as a unified sample with identified blazars outside the mask.
 Tn this case. as diseussed in 8??.. the flax function slope is considerablv shallower. allowing for a more extended extrapolation to low fluxes before the ECRB is overproduced.," In this case, as discussed in \ref{cfd}, the flux function slope is considerably shallower, allowing for a more extended extrapolation to low fluxes before the EGRB is overproduced."
 The lowest flux limit allowed for the extrapolation iu this case is shown with the erey line in refloguloes.. aud 1s amoderate factor ~3 below the limits of the unidentified source samples when considered alouc.," The lowest flux limit allowed for the extrapolation in this case is shown with the grey line in \\ref{lognlogs}, and is a moderate factor $\sim 3$ below the limits of the unidentified source samples when considered alone."
 A comparison between the of the cumulative Cluission spectrum of unresolved unidentified sources aud 16 EGRET EGRB is shown in refress.. The dashed line shows a sinele power-law fit to 1ο Sreckuimar et ((1998) determination of the ECRB.," A comparison between the of the cumulative emission spectrum of unresolved unidentified sources and the EGRET EGRB is shown in \\ref{ress}, The dashed line shows a single power-law fit to the Sreekumar et (1998) determination of the EGRB."
 The thin solid liue is the more recent redetermination of je exannun-rayv backeround by Stroug et ((200la). in which they used a more detailed model of the NBlkvy Wavy diffuse emission to subtract the Galactic componout rom the EGRET diffuse sky map.," The thin solid line is the more recent redetermination of the gamma-ray background by Strong et (2004a), in which they used a more detailed model of the Milky Way diffuse emission to subtract the Galactic component from the EGRET diffuse sky map."
 The dotted lines are he systematic uncertaimties in the ECRB determination of Strong et 1a). entering through their model of the Galaxy.," The dotted lines are the systematic uncertainties in the EGRB determination of Strong et (2004a), entering through their model of the Galaxy."
 Our ((200calculation of the spectrum of the nuresolved unidentified source component is shown witli the thick solid line., Our calculation of the spectrum of the unresolved unidentified source component is shown with the thick solid line.
 The iiaxiuunc-likelihood parameters ποια Sample 1 were used for pla) in this calculation. although. as becomes clear from refcout in the Appendix. the spectral shapes resulting from all three samples are consistent with each other.," The maximum-likelihood parameters from Sample 1 were used for $p(\alpha)$ in this calculation, although, as becomes clear from \\ref{cont} in the Appendix, the spectral shapes resulting from all three samples are consistent with each other."
 At low energies. where the systematics are low. the spectrmm of the unideutiied. component is in excellent agreement with the ECGRD observational spectrum of Strong et ((2001a4).," At low energies, where the systematics are low, the spectrum of the unidentified component is in excellent agreement with the EGRB observational spectrum of Strong et (2004a)."
 At higher energies. where je systematics are large. the unideutified coniponeut spectruni is hugelv within systematics except at very Heh energies.," At higher energies, where the systematics are large, the unidentified component spectrum is largely within systematics except at very high energies."
 If unidentified sources are indeed a lominant contribution at relatively low energies. then us result inav be perceived as a tantalizing hint that at 1ο highest enereics of the ECRET range a new type of contribution. (c.¢.. from ligh-enerey peaked BL Lacs. or roni anunihilatius dark matter) may become miportaut at ~L1 GeV.," If unidentified sources are indeed a dominant contribution at relatively low energies, then this result may be perceived as a tantalizing hint that at the highest energies of the EGRET range a new type of contribution, (e.g., from high-energy peaked BL Lacs, or from annihilating dark matter) may become important at $\sim 1{\rm \, GeV}$ ."
 It should be noted that the amplitude of the cumulative spectra plotted in does uot constitute a prediction: rather. itis," It should be noted that the amplitude of the cumulative spectrum plotted in \\ref{ress} does not constitute a prediction; rather, itis"
ihe moment.,the moment.
 The second factor should be maximizecl., The second factor should be maximized.
" Since the readout time is fixed. the smaller is the telescope aperture. the longer can be the exposures before a Vii,=8 mag star saturates. and so the smaller fraction of time is lost to read-out."," Since the readout time is fixed, the smaller is the telescope aperture, the longer can be the exposures before a $V_{\rm min}=8$ mag star saturates, and so the smaller fraction of time is lost to read-out."
" For large-format detectors. the pixel size is (vpically Ar,=15jim. for which significant non-linearities set in at about 60.000€."," For large-format detectors, the pixel size is typically $\Delta x_p=15\,\mu$ m, for which significant non-linearities set in at about $60,000\,e^-$."
" lence. exposure “mes are That is. the exposure time for à 1"" ""telescope"" is already of the order of typical read-out times for large-Iormat. detectors."," Hence, exposure times are That is, the exposure time for a $1''$ “telescope” is already of the order of typical read-out times for large-format detectors."
 Clearly. smaller is better. but are there constraints [rom eoimg (oo small?," Clearly, smaller is better, but are there constraints from going too small?"
 One potential constraint comes Iron sky noise., One potential constraint comes from sky noise.
 To stay within the photon-noise limited reeime. which has been assumed in all of our calculations. the skv inside one pixel should be at least one magnitude fainter (han Vy. and preferably two mags fainter.," To stay within the photon-noise limited regime, which has been assumed in all of our calculations, the sky inside one pixel should be at least one magnitude fainter than $V_{\rm max}$, and preferably two mags fainter."
" Assuming the sky brightness in our broadening passband reaches a maximum of V.—19.7arcsec.7. and again assuming Ar,=15 an pixels. the sky in one pixel is ]lence. the skv is a bit bright for a 1 telescope. but appears quite satisfactory for a 2""."," Assuming the sky brightness in our broadening passband reaches a maximum of $V=19.7\,\rm arcsec^{-2}$, and again assuming $\Delta x_p=15\,\mu$ m pixels, the sky in one pixel is Hence, the sky is a bit bright for a $1''$ telescope, but appears quite satisfactory for a $2''$."
 Finally. one must be careful that the field of view is not too large. or the local-plane distortions at its edges will be difficult (and expensive) to correct.," Finally, one must be careful that the field of view is not too large, or the focal-plane distortions at its edges will be difficult (and expensive) to correct."
" For example. for a thx4h detector with Ar,=15 yan pixels. M9=687(D/2.5em)'(F/2)|."," For example, for a $4k\times 4k$ detector with $\Delta x_p=15\,\mu$ m pixels, $\Delta \theta = 68^\circ (D/2.5\,{\rm cm})^{-1}({\cal F}/2)^{-1}$."
" ACI"". this is probably too large to correct al reasonable expense."," At $1''$, this is probably too large to correct at reasonable expense."
" With tliis size detector. both skv-noise considerations. and problems in opties design argue for a 1.5"" telescope."," With this size detector, both sky-noise considerations, and problems in optics design argue for a $1.5''$ telescope."
" However. from equation (6)). such a ""Jarge"" telescope will most likely be dominated by read-out time."," However, from equation \ref{equtexp}) ), such a “large” telescope will most likely be dominated by read-out time."
" In summary. we conclude that al” to 2"" telescope equipped with a tax4h. £= Gem detector and with F=2 local ralio is optimal for this observing problem."," In summary, we conclude that a $1''$ to $2''$ telescope equipped with a $4x\times 4k$, ${\cal L}=6\,$ cm detector and with ${\cal F}=2$ focal ratio is optimal for this observing problem."
" Among all existing transit programs of which we are aware. the WASP telescope (2=2.5"" lens. F=2.8 focal ratio. 2hx 2k. £= 3em detector. Streetetal. 2002)) comes closest to meeting (hese design specifications."," Among all existing transit programs of which we are aware, the WASP telescope $D=2.5''$ lens, ${\cal F}=2.8$ focal ratio, $2k \times 2k$ , ${\cal L}=3\,$ cm detector, \citealt{str02}) ) comes closest to meeting these design specifications."
 Adopting +—u. AK=Ny. ἕν=20%. ὃς=50. ο= Gem. and F=2. we find from equation (5)) that (he required duration of the experiment using our optimallv-desiened telescope is J=3inonths.," Adopting $\gamma=\gamma_0$, $K=K_0$, ${\cal E}_{0}=20\%$, ${\cal E}_{S}=50\%$, ${\cal L}=6\,$ cm, and ${\cal F}=2$, we find from equation \ref{equGam2}) ) that the required duration of the experiment using our optimally-designed telescope is $T=3\,\rm months$."
 Clearly. roughly à vear is required just to get around (he sky.," Clearly, roughly a year is required just to get around the sky."
 The S/N acquired. during such a vear-long search would therefore roughly double the minimum requirements calculated here., The S/N acquired during such a year-long search would therefore roughly double the minimum requirements calculated here.
 This work was supported in part bv grant. AST 02-01266 [rom theNSF., This work was supported in part by grant AST 02-01266 from theNSF.
Tl16 predicted frequencies discussed above depend oulv on the observed frequencies.,The predicted frequencies discussed above depend only on the observed frequencies.
" It the observed frequeucies are rj aL Myon. then the predicted frequency in the relativistic precession model is 254,4Moy and the predicted frequency in the beat frequency model is δν."," If the observed frequencies are $\nu_{\rm low}$ and $\nu_{\rm high}$, then the predicted frequency in the relativistic precession model is $2\nu_{\rm high}-\nu_{\rm low}$ and the predicted frequency in the beat frequency model is $2\nu_{\rm low}$."
 The frequencies will change with time for a giveu ΝΟος but these predictions are ideally suited for testing by use of the method first used by Méuudez ((1998) to ¢iscover a weak brightuess oscillation iu. LU 160852.," The frequencies will change with time for a given source, but these predictions are ideally suited for testing by use of the shift-and-add method first used by Ménndez (1998) to discover a weak brightness oscillation in 4U 1608–52."
 The kmuch of anew ecueration of high spectral resolution satellites such as Claudra and NATAL opens up new wavs to probe the strong eravity and deuse matter of jeutron stars., The launch of a new generation of high spectral resolution satellites such as Chandra and XMM opens up new ways to probe the strong gravity and dense matter of neutron stars.
 For example. suppose that simu]taneous observations of a kilohertz QPO source are performed with RATE and NADL.," For example, suppose that simultaneous observations of a kilohertz QPO source are performed with and XMM."
 Π an Fe Ίνα profile from the inner edge of the nearly-circular flow is detected and characterized with NMM. this xofile gives the radius of that inuer edge iu uuts of the eravitational mass of the jeutron star.," If an Fe $\alpha$ profile from the inner edge of the nearly-circular flow is detected and characterized with XMM, this profile gives the radius of that inner edge in units of the gravitational mass of the neutron star."
 If at the same time a pair of kilorertz QPOs is detected withZXTE. ieu the frequency of the upper peals. combired with the radius in uuits of the eravitational mass. viclds the mass of the star.," If at the same time a pair of kilohertz QPOs is detected with, then the frequency of the upper peak, combined with the radius in units of the gravitational mass, yields the mass of the star."
 This is true independent of where ie inner edge is: that is. it need not be at he innermost stable circular orbit.," This is true independent of where the inner edge is; that is, it need not be at the innermost stable circular orbit."
 Therefore. such simultaneous measurements could provide a clean way to mcasure je eravitational masses of neutron stars m low-niass X-ray binaries.," Therefore, such simultaneous measurements could provide a clean way to measure the gravitational masses of neutron stars in low-mass X-ray binaries."
 Future. ligh-area missions such as Coustellation-X. or a livpothetical high-arca ollow-up tiniic ussion to could provide even more information. by allowing tharacterization of the waveforms of the brieltuess oscillations seen ia accretiou-oowered and burst-powered emission.," Future, high-area missions such as Constellation-X or a hypothetical high-area follow-up timing mission to could provide even more information, by allowing characterization of the waveforms of the brightness oscillations seen in accretion-powered and burst-powered emission."
" Au example is shown in the left panel of Figure 2. which show"" theoretical waveforms for burst briehtuess oscillations."," An example is shown in the left panel of Figure 2, which shows theoretical waveforms for burst brightness oscillations."
 The curve for the more compact star is broader. as is expected due to the extra gravitational ight deflection.," The curve for the more compact star is broader, as is expected due to the extra gravitational light deflection."
 In addition. the curve for the larger (less compact) star is more asvluuetric. due to the ereater Doppler shifts from 1ο surface rotation velocity.," In addition, the curve for the larger (less compact) star is more asymmetric, due to the greater Doppler shifts from the surface rotation velocity."
 The waveform therefore encodes information about |oth the mass and radius of he sar. nicaning that repeated observations of bursts from a single source can constrain the mass aud radius tightly. with cousequent coustraints on the equation of state of the high-density matter in the core of neutro) stars.," The waveform therefore encodes information about both the mass and radius of the star, meaning that repeated observations of bursts from a single source can constrain the mass and radius tightly, with consequent constraints on the equation of state of the high-density matter in the core of neutron stars."
 Tie nieht panel of Figure 2 shows an exae of tlie' constraints possible with Cosellation-N. (outer contour. at 16) aud a hypothBical 10 i10? future timing institment Guner contour. at lo).," The right panel of Figure 2 shows an example of the constraints possible with Constellation-X (outer contour, at $\sigma$ ) and a hypothetical 10 $^2$ future timing instrument (inner contour, at $\sigma$ )."
 Clearly. waveform fittine can in principle vield verv xecise Information about the mass auk racius of ilividual neutron stars. aud thereore about the equation of state of mater at supramiclear densities.," Clearly, waveform fitting can in principle yield very precise information about the mass and radius of individual neutron stars, and therefore about the equation of state of matter at supranuclear densities."
 Ilieh-area timing nissious also will poteutialy allow us to do qualitatively new tvpes of tests of the stroug-eravitv predictiois of eoncral relativity., High-area timing missions also will potentially allow us to do qualitatively new types of tests of the strong-gravity predictions of general relativity.
 For example. we have just discussed two iudependoeut ways of ostnuatius he gravitational mass of a neutron star: bv combining Fe Ίνα profiles with kilohertz QPOs. aud by waveform fitting of burst brightuess oscillations.," For example, we have just discussed two independent ways of estimating the gravitational mass of a neutron star: by combining Fe $\alpha$ profiles with kilohertz QPOs, and by waveform fitting of burst brightness oscillations."
 If either o these is successful for au iudividual source. then we can predict preciscly the orvital frequency at the imuermost stable circular orbit.," If either of these is successful for an individual source, then we can predict precisely the orbital frequency at the innermost stable circular orbit."
 With ligh-areatimine mstruneuts we expect to see more cases, With high-areatiming instruments we expect to see more cases
angular spectra for a coupled cosmology. in which dark matter decays into dark energy.,"angular spectra for a coupled cosmology, in which dark matter decays into dark energy."
"Moreover, the acceleration vector (of EC 95b) found in our fit appears to point almost exactly towards the expected position of EC 95a at the median epoch of our observations 55a).","Moreover, the acceleration vector (of EC 95b) found in our fit appears to point almost exactly towards the expected position of EC 95a at the median epoch of our observations 5a)."
 This is exactly what would be expected if EC Όσα and EC 95b formed a gravitauonally bound binary system., This is exactly what would be expected if EC 95a and EC 95b formed a gravitationally bound binary system.
" Finally, since the (trajectory of EC 95b is curved and accelerated, whereas that of EC 95a is linear and uniform, EC 95a must be significantly more massive than EC 95b."," Finally, since the trajectory of EC 95b is curved and accelerated, whereas that of EC 95a is linear and uniform, EC 95a must be significantly more massive than EC 95b."
" We conclude that EC 95a and EC 95b do constitute a binary system, where EC 95a is the primary and EC 95b the secondary."," We conclude that EC 95a and EC 95b do constitute a binary system, where EC 95a is the primary and EC 95b the secondary."
" This is, indeed, the reason why we ascribed those names to the sources in the first place."," This is, indeed, the reason why we ascribed those names to the sources in the first place."
" To further characterize the system, it 15 useful to calculate the separation between EC 95 a and EC 95b as a l'unction of me."," To further characterize the system, it is useful to calculate the separation between EC 95 a and EC 95b as a function of time."
" Since the two sources are detected simultaneously at only two epochs, this can be unambiguously done only at those two epochs."," Since the two sources are detected simultaneously at only two epochs, this can be unambiguously done only at those two epochs."
" However, since the motion of EC 95a appears to be very nearly linear and uniform, we can estimate the position of the primary at the other five epochs when the secondary is detected."," However, since the motion of EC 95a appears to be very nearly linear and uniform, we can estimate the position of the primary at the other five epochs when the secondary is detected."
 Those positions are shown as blue solid circles in 55a., Those positions are shown as blue solid circles in 5a.
" For epochs 6 and 7, the estimation of the position of the primary only involves an interpolation."," For epochs 6 and 7, the estimation of the position of the primary only involves an interpolation."
" However, for the first three epochs, somewhat more uncertain extrapolauons are needed."," However, for the first three epochs, somewhat more uncertain extrapolations are needed."
" Using these esumates [or the position of the primary, we calculated the separation between the two members of the system shown in 55b."," Using these estimates for the position of the primary, we calculated the separation between the two members of the system shown in 5b."
" It is clear from that figure that our observations only cover a fairly small fraction of the orbit, and that any orbit modeling will be very uncertain."," It is clear from that figure that our observations only cover a fairly small fraction of the orbit, and that any orbit modeling will be very uncertain."
" A poorly constrained, but plausible fit (see below) is shown as a solid curve."," A poorly constrained, but plausible fit (see below) is shown as a solid curve."
 The dotted curve corresponds to the parabolic path predicted by our uniformly accelerated fit., The dotted curve corresponds to the parabolic path predicted by our uniformly accelerated fit.
 The fact that those two wrajectories are nearly indistnguishable over the course of our observations justifies a posteriori the use of a uniformly accelerated motion in our astrometric fit., The fact that those two trajectories are nearly indistinguishable over the course of our observations justifies a posteriori the use of a uniformly accelerated motion in our astrometric fit.
 A strict lower limit to the mass of the primary can be found from the magnitude of the acceleration vector determined earher., A strict lower limit to the mass of the primary can be found from the magnitude of the acceleration vector determined earlier.
" II 7: and M are the masses of the secondary and the primary, respectively, Newton's law applied to the secondary yields: where « is the magnitude of the acceleration (of the secondary), and 7 is the true separation between the sources."," If $m$ and $M$ are the masses of the secondary and the primary, respectively, Newton's law applied to the secondary yields: where $a$ is the magnitude of the acceleration (of the secondary), and $r$ is the true separation between the sources."
" Of course, we measure only quantities projected onto the plane of the sky, so the measured separation and acceleration are only lower limits."," Of course, we measure only quantities projected onto the plane of the sky, so the measured separation and acceleration are only lower limits."
" From the measured values (o,,5, =2.4 mas yr ?m (015 cm ? and fj, = 15.8 mas = 9.86 x 10 em), we obtain à minimum mass for EC 95a of 1.1 M.."," From the measured values $a_{min}$ =2.4 mas $^{-2}$ $\equiv$ 0.015 cm $^{-2}$ and $r_{min}$ = 15.8 mas $\equiv$ 9.86 $\times$ $^{13}$ cm), we obtain a minimum mass for EC 95a of 1.1 $M_\odot$."
" The mass goes roughly as cos?i, so if the inclination were 45"", the mass of EC 95a would be about 3 M..."," The mass goes roughly as $\cos^3{i}$, so if the inclination were $^\circ$, the mass of EC 95a would be about 3 $M_\odot$."
" To obtain an alternative mass estimate, we model the component relative astrometry data 11: 55b) with a Keplerian orbit."," To obtain an alternative mass estimate, we model the component relative astrometry data 1; 5b) with a Keplerian orbit."
" As these data apparently cover only a small (~ 10%)) Iraction of the EC 95 orbit, we constrained the orbit model to be circular (e = 0); this assumption is not physically motivated, but the data do not presently support a more complex interpretation."," As these data apparently cover only a small $\sim$ ) fraction of the EC 95 orbit, we constrained the orbit model to be circular $e$ = 0); this assumption is not physically motivated, but the data do not presently support a more complex interpretation."
 The result of our modeling is shown as a solid line in Fig., The result of our modeling is shown as a solid line in Fig.
 5b., 5b.
" The semi-major axis of this preliminary orbit is 31 + 9 mas, its inclination is 7 = 60° + 10°, and the orbital period is 16.5 5.0 yr."," The semi-major axis of this preliminary orbit is 31 $\pm$ 9 mas, its inclination is $i$ = $^\circ$ $\pm$ $^\circ$, and the orbital period is 16.5 $\pm$ 5.0 yr."
 These orbit parameters imply à total mass for the system of M=8'!&27 M.., These orbit parameters imply a total mass for the system of $M$ = $^{+27}_{-6}$ $M_\odot$ .
" The large error on the upper limit⋅⋅ reflects the fact. that the mass goes roughly as cosiei, and"," The large error on the upper limit reflects the fact that the mass goes roughly as $\cos^3{i}$, and"
"analyses in ? and ?,, where the total dark matter velocity dispersion was inferred assuming either 6=0, or the analytical 6-profiles of ? or? (see also ?)).","analyses in \citet{2004ApJ...611..175I} and \citet{2007MNRAS.380.1521M}, where the total dark matter velocity dispersion was inferred assuming either $\beta=0$, or the analytical $\beta$ -profiles of \citet{2000ApJ...539..561C} or \citet{1996MNRAS.281..716C} (see also \citet{2008MNRAS.tmp..719W}) )."
" In particular, ? found that the dark matter temperature and the ICM temperature were essentially the same in strong cooling-core clusters."," In particular, \citet{2007MNRAS.380.1521M} found that the dark matter temperature and the ICM temperature were essentially the same in strong cooling-core clusters."
" 'The structure of the paper is the following: In the next section, we discuss how we relate the temperature of dark matter to the observable gas temperature."," The structure of the paper is the following: In the next section, we discuss how we relate the temperature of dark matter to the observable gas temperature."
 In section 3 we show how we can then solve the dynamics of the dark matter., In section 3 we show how we can then solve the dynamics of the dark matter.
" In section 4 we test the assumed temperature relation and our method on numerical simulations, and in section 5 we apply the method to observational data."," In section 4 we test the assumed temperature relation and our method on numerical simulations, and in section 5 we apply the method to observational data."
 Section 6 is the summary and discussion., Section 6 is the summary and discussion.
 'The equality of inertial and gravitational mass implies that the orbit of a test particle in a gravitational system is independent of mass., The equality of inertial and gravitational mass implies that the orbit of a test particle in a gravitational system is independent of mass.
" For example, the velocity of a circular orbit in a spherical mass distribution v2—GM depends only on the distance to the center of the system(r)/r and the mass contained within that radius."," For example, the velocity of a circular orbit in a spherical mass distribution $v_c^2=GM(r)/r$ depends only on the distance to the center of the system and the mass contained within that radius."
" Therefore it is natural to assume that, at a given radius, all species in a relaxed, spherical gravitational system have the same average specific kinetic energy."," Therefore it is natural to assume that, at a given radius, all species in a relaxed, spherical gravitational system have the same average specific kinetic energy."
" Obviously, they also have the same specific potential energy."," Obviously, they also have the same specific potential energy."
" In a gas system, equilibrium implies energy equipartition between all species."," In a gas system, equilibrium implies energy equipartition between all species."
" It is clear that the corresponding principle for a relaxed gravitational system is a common velocity dispersion, precisely because gravitational dynamics are independent of mass."," It is clear that the corresponding principle for a relaxed gravitational system is a common velocity dispersion, precisely because gravitational dynamics are independent of mass."
" Since the average velocity is associated with the thermal energy content, this relationship is expressed by The parameter & is constant as long as the impact of radiative or entropy-changing processes affecting the gas is negligible and the system is relaxed."," Since the average velocity is associated with the thermal energy content, this relationship is expressed by The parameter $\kappa$ is constant as long as the impact of radiative or entropy-changing processes affecting the gas is negligible and the system is relaxed."
" Therefore, we allow for a radial dependence, κ.=&(r/rosoo), where T2500 is the scale radius within which the mean total density is 2500 times the critical density at the redshift of the cluster."," Therefore, we allow for a radial dependence, $\kappa=\kappa (r/r_{2500})$, where $r_{2500}$ is the scale radius within which the mean total density is $2500$ times the critical density at the redshift of the cluster."
 The dark matter temperature in is naturally not well-defined as there is no thermodynamic(2)) equilibrium for a collisionless gas., The dark matter temperature in \ref{eq:trl}) ) is naturally not well-defined as there is no thermodynamic equilibrium for a collisionless gas.
" Instead, we simply define an effective dark matter temperature which is proportional to the three-dimensional velocity dispersion, The velocity dispersion has been decomposed into the contributions from the one-dimensional radial and tangential dispersions."," Instead, we simply define an effective dark matter temperature which is proportional to the three-dimensional velocity dispersion, The velocity dispersion has been decomposed into the contributions from the one-dimensional radial and tangential dispersions."
 We choose the constant of proportionality to be the mean molecular mass of the intracluster gas simply to allow & to be of order unity., We choose the constant of proportionality to be the mean molecular mass of the intracluster gas simply to allow $\kappa$ to be of order unity.
" are equivalent to assuming that the specific Equationsenergies of (2))-@))gas and dark matter particles are the same up to a factor of &, on average."," Equations \ref{eq:trl}) \ref{eq:tdm}) ) are equivalent to assuming that the specific energies of gas and dark matter particles are the same up to a factor of $\kappa$, on average."
 The same conjecture was made in 7? but with «=1 explicitly., The same conjecture was made in \citet{2007A&A...476L..37H} but with $\kappa=1$ explicitly.
 It should be mentioned that the temperature relation was recently analyzed in simulations by ?.., It should be mentioned that the temperature relation \ref{eq:trl}) ) was recently analyzed in simulations by \citet{2008ApJ...672..122E}.
" Whereas we (9)allow a possible radial variation in the temperature relation, those authors considered averages within rooo and found that This was based on their determination of =+0.04)(Typec/Tmw), where the ratio ofΕλ the spectroscopic(0.87 temperature to the mass-weighted temperature was =1.1+0.05 (?).."," Whereas we allow a possible radial variation in the temperature relation, those authors considered averages within $r_{200}$ and found that This was based on their determination of $\bar{\kappa}^{-1}_{<r_{200}}=(0.87\pm0.04)\langle T_\mathrm{spec}/T_\mathrm{mw}\rangle$, where the ratio of the spectroscopic temperature to the mass-weighted temperature was $\langle T_\mathrm{spec}/T_\mathrm{mw}\rangle=1.1\pm0.05$ \citep{2006ApJ...650..538N}."
" Then, in ? it was noted(Tspec/Tmw) that by applying virial scaling to the WMAP5+SN+BAO results (?),, an average value of &=1.1 was found."," Then, in \citet{2008ApJ...679L...1R} it was noted that by applying virial scaling to the WMAP5+SN+BAO results \citep{2008arXiv0803.0547K}, an average value of $\tilde{\kappa}^{-1}|_{<r_{500}}=1.1$ was found."
" The authors concluded that l|5,,,the observational results indicated that the average specific energy of the ICM was close to both that of the dark matter and that of the galaxies.", The authors concluded that the observational results indicated that the average specific energy of the ICM was close to both that of the dark matter and that of the galaxies.
" In section ??,, we will arrive at the same conclusion for simulated galaxy clusters."," In section \ref{sec:sim}, we will arrive at the same conclusion for simulated galaxy clusters."
 Equation (2)) allows us to estimate the total velocity dispersion profile of the dark matter structure from measurements of the radial temperature profile of the gas., Equation \ref{eq:trl}) ) allows us to estimate the total velocity dispersion profile of the dark matter structure from measurements of the radial temperature profile of the gas.
 In this section we discuss how we can proceed to determine the dark matter velocity anisotropy., In this section we discuss how we can proceed to determine the dark matter velocity anisotropy.
 The collisionless Jeans equations relate the dynamical properties of the dark matter to the gravitational potential of the cluster., The collisionless Jeans equations relate the dynamical properties of the dark matter to the gravitational potential of the cluster.
" Assuming that the system is spherically symmetric and in a steady state, the second Jeans equation can be put in the form (?) where v is the dark matter number density, v? is the second moment of the ith velocity component, and M is the mass contained within radiusr."," Assuming that the system is spherically symmetric and in a steady state, the second Jeans equation can be put in the form \citep{1987gady.book.....B}
 where $\nu$ is the dark matter number density, $\overline{v_i^2}$ is the second moment of the $i$ th velocity component, and $M$ is the mass contained within radius$r$."
" If it is further assumed that there are no bulk flows, 0;=0, and that the tangential velocity dispersions are equal, 05=0707, the Jeans equation becomes where ppy is the mass density, is the radial velocity dispersion, and f is the velocity anisotropyc2 introduced in (1))."," If it is further assumed that there are no bulk flows, $\overline{v_i}=0$, and that the tangential velocity dispersions are equal, $\sigma_\theta^2=\sigma_\phi^2\equiv\sigma_t^2$, the Jeans equation becomes where $\rho_\mathrm{DM}$ is the mass density, $\sigma_r^2$ is the radial velocity dispersion, and $\beta$ is the velocity anisotropy introduced in \ref{eq:be}) )."
" (D)lar to the Jeans equation, the radial part of the Euler equations of the ICM expresses the condition that the thermal pressure of the gas balances the gravitational potential."," Similar to the Jeans equation, the radial part of the Euler equations of the ICM expresses the condition that the thermal pressure of the gas balances the gravitational potential."
" This equation of hydrostatic equilibrium reads where is the gas electron temperature and n, is the number densityTyas of electrons.", This equation of hydrostatic equilibrium reads where $T_\mathrm{gas}$ is the gas electron temperature and $n_e$ is the number density of electrons.
 This important equation has been widely used to estimate the total mass of a galaxy cluster from X-ray data., This important equation has been widely used to estimate the total mass of a galaxy cluster from X-ray data.
 In case there is turbulence or larger scale bulk motion in the gas additional terms of the form (¢-V)U—(vj+ appear (?).., In case there is turbulence or larger scale bulk motion in the gas additional terms of the form $(\vec{v}\cdot\nabla)\vec{v}-(v_\theta^2+v_\phi^2)/r$ appear \citep{lanlif}. .
" Neglecting such terms may lead to an v3)/runderestimate of the mass, however this is usually not a major effect for systems that appear relaxed (?).."," Neglecting such terms may lead to an underestimate of the mass, however this is usually not a major effect for systems that appear relaxed \citep{2008arXiv0808.1111P}. ."
Figures 6 and 7 display the mean light curves folded on ; using the 2002 and 2003 data set. and Po using the 2003 data set. respectively.,"Figures 6 and 7 display the mean light curves folded on $_1$ using the 2002 and 2003 data set, and $_2$ using the 2003 data set, respectively."
" Ehe shape of the profiles are sinusoidal or both P, and P» (in 2002-2003) with a slight asymmetry.", The shape of the profiles are sinusoidal for both $_1$ and $_2$ (in 2002-2003) with a slight asymmetry.
 Llump-like features are apparent on both profiles. as well.," Hump-like features are apparent on both profiles, as well."
" ‘The ingress in the mean light curve of P, and P» is shorter w LOM of the photometric phase than the egress.", The ingress in the mean light curve of $_1$ and $_2$ is shorter by $\%$ of the photometric phase than the egress.
 In adition. we have constructed. color diagrams VHR. and Il.V to search for color variations over the photometric yhases of the periods (light curves are calibrated before the construction of the color diagrams).," In adition, we have constructed color diagrams $V-R$, $I-R$, and $I-V$ to search for color variations over the photometric phases of the periods (light curves are calibrated before the construction of the color diagrams)."
 The right-hand panel of Fig., The right-hand panel of Fig.
 S shows the existence of a significant modulation of he color dillerence αςV over the photometric phase of Py., 8 shows the existence of a significant modulation of the color difference $I-V$ over the photometric phase of $_1$.
 The fl [magnitudes plotted on the left-hand panel of Eig., The $I-R$ magnitudes plotted on the left-hand panel of Fig.
 S indicates color variation over the photometric phase of the second period as well., 8 indicates color variation over the photometric phase of the second period as well.
 We presented Lt nights of data on V2275 (νο obtained with the PUG 1.5 m telescope using mainly standard: 7? filters in 2002 and 2003., We presented 14 nights of data on V2275 Cyg obtained with the TUG 1.5 m telescope using mainly standard $R$ filters in 2002 and 2003.
" We discover large mocdulations Am,=0.42 at D,20.31449(15) d in the light curve of the classical nova in 2002 and the amplitude of this modulation decrease to. Am,=0.22 in 2003.", We discover large modulations $\Delta m_r$ =0.42 at $_1$ =0.31449(15) d in the light curve of the classical nova in 2002 and the amplitude of this modulation decrease to $\Delta m_r$ =0.22 in 2003.
 This period. is discovered in the V. and the £ band as well., This period is discovered in the $V$ and the $I$ band as well.
 Since the periodicity is persistent and. coherent. we propose that this is the binary period. of the svstem.," Since the periodicity is persistent and coherent, we propose that this is the binary period of the system."
 The orbital period. can be detected as a result of the aspect. variations of the secondary. due to heating from the hot WD (Ixovetz. Prialnik. & Shara LOSS).," The orbital period can be detected as a result of the aspect variations of the secondary due to heating from the hot WD (Kovetz, Prialnik, $\&$ Shara 1988)."
 The decrease in the modulation depth from 2002 to 2003 and the color variations (4 1.1. RORH 1) support a scenario where the heated secondary is the source of the light and the color variation rather than à hot spot on the disk.," The decrease in the modulation depth from 2002 to 2003 and the color variations $I-V$, $I-R$, $R-V$ ) support a scenario where the heated secondary is the source of the light and the color variation rather than a hot spot on the disk."
" “Phe Alass-Period relation Mo2:0.11D,,. (Warner 1995) vields a secondary mass of about 0.83 M. lor V2275 (νο using the discovered. binary period in this paper. which is in accordance with the fact that a massive secondary will be hot and prone to irradiation effects."," The Mass-Period relation $_2$$\simeq$ $_{hr}$ (Warner 1995) yields a secondary mass of about 0.83 $_{\odot}$ for V2275 Cyg using the discovered binary period in this paper, which is in accordance with the fact that a massive secondary will be hot and prone to irradiation effects."
 In. addition. the mean light. curves (Fig.," In addition, the mean light curves (Fig."
 6 and 7) indicate an asymmetry in the 1t band as would be expected. from the dillerential, 6 and 7) indicate an asymmetry in the R band as would be expected from the differential
reach (he shot noise limit for ordinary spectra with integration times at least as deep as 800.000 sec.,"reach the shot noise limit for ordinary spectra with integration times at least as deep as 800,000 sec."
 We treat the shot noise in an 800.000 sec integration as a limiting uncertainty for deeper observations. not because of evidence that. svstematic effects become significant al longer times but because we have not vet shown their absence.," We treat the shot noise in an 800,000 sec integration as a limiting uncertainty for deeper observations, not because of evidence that systematic effects become significant at longer times but because we have not yet shown their absence."
 The wavelength seale is based on preflight measurements. offset bv a (constant) 1 determined from the bright Ile 256.3 feature. which is present in essentially all the observations.," The wavelength scale is based on preflight measurements, offset by a (constant) $\sim$ 1 determined from the bright He 256.3 feature, which is present in essentially all the observations."
 Temporal variations in the measured centroid of the Ile feature shows a dispersion of about 0.2/., Temporal variations in the measured centroid of the He feature shows a dispersion of about 0.2.
. Periodic observations of the moon. which provides a weakly reflected. solar spectrum (hat includes (he features of Fe (through.NIT. confirm that the adopted wavelength scale provides a good fit near the center of the spectral band. aud (that the relative throughput of the filler panels is as measured pre-flight.," Periodic observations of the moon, which provides a weakly reflected solar spectrum that includes the features of Fe through, confirm that the adopted wavelength scale provides a good fit near the center of the spectral band, and that the relative throughput of the filter panels is as measured pre-flight."
 The measured fIuxes of both the He feature and the Iunar spectrum are in good agreement with pre-Ilight expectations. suggesting (hat ihe instrument throughput has not declined since laboratory calibration.," The measured fluxes of both the He feature and the lunar spectrum are in good agreement with pre-flight expectations, suggesting that the instrument throughput has not declined since laboratory calibration."
 Additional details of the on-orbit verification will be presented in a subsequent work., Additional details of the on-orbit verification will be presented in a subsequent work.
 Over most of the CLIPS spectral band. and certainly near the wavelengths of the iron lines. photoelectric absorption from (he neutral interstellar medium limits the observable path length (to al most a few hundred parsecs.," Over most of the CHIPS spectral band, and certainly near the wavelengths of the iron lines, photoelectric absorption from the neutral interstellar medium limits the observable path length to at most a few hundred parsecs."
 Thus the only astroplivsical source we expect lo see is emission from the local hot bubble that should be isotropic (within a factor of a few) on the sky., Thus the only astrophysical source we expect to see is emission from the local hot bubble that should be isotropic (within a factor of a few) on the sky.
 In Figure 2 we show a spectrum from 150 - 200/.. binned at 0.5/.," In Figure 2 we show a spectrum from 150 - 200, binned at 0.5."
". This is the complete spectrum lor the CY 2003-2004 narrow slit program. repesenting 13.2 Msec of effective observing (ime. inclusive of all latencies or ""dead times.”"," This is the complete spectrum for the CY 2003-2004 narrow slit program, repesenting 13.2 Msec of effective observing time, inclusive of all latencies or ""dead times."""
 About of the data were collected during orbital dav: during orbital night., About of the data were collected during orbital day; during orbital night.
 Both filter panels ave included in the histogram., Both filter panels are included in the histogram.
 Much of the observing (ime was directed al regions where (he EUV brightness was expected to be hieh based on high soft X-ray brightness. low column density of neutral hydrogen. ete.," Much of the observing time was directed at regions where the EUV brightness was expected to be high based on high soft X-ray brightness, low column density of neutral hydrogen, etc."
 In Figure 2 we also show a scaled. charged-particle background spectrum and an APEC/CIILANTI model (discussed below) with parameters [rom one model lor the local hot bubble (emission measure 0.0038 ° pc. log(D)—6.1.," In Figure 2 we also show a scaled, charged-particle background spectrum and an APEC/CHIANTI model (discussed below) with parameters from one model for the local hot bubble (emission measure 0.0038 $^{-6}$ pc, log(T)=6.1,"
"(3j= (O74. fo> 35%) is in the beginning and LED. 141569 (Ji—Ll.T4. f,= LA) is in the end of the sequence.","$\beta_1= +0.74$ , $f_c > 35\%$ ) is in the beginning and HD 141569 $\beta_1 = -1.74$, $f_c = 1.4\%$ ) is in the end of the sequence."
 The diminishing of { with the SED slope is in agreement with he ILXeDe evolutionary sequence suggested by Malfaitοἱal. (1998)., The diminishing of $f_c$ with the SED slope is in agreement with the HAeBe evolutionary sequence suggested by \cite{Malfait98}.
. In à comparison with the circumstellar. geometries. of he two major LLAeBes groups proposed by Meeus (2002).. the low f. of LID 141549 is consistent with their suggestion of a flat disc structure for Croup EH. while the intermediary. f. values for BD-14 1319 and LAS 06475-0735 are consistent with the optically thick disc surrounded by a Ilared. thin region. as proposed. for their Group Ll. On the extreme is HAS 07394-1953. showing the highest f. of the sample. which precedes the two groups of τους et al.," In a comparison with the circumstellar geometries of the two major HAeBes groups proposed by \cite{Meeus02}, the low $f_c$ of HD 141549 is consistent with their suggestion of a flat disc structure for Group II, while the intermediary $f_c$ values for BD-14 1319 and IRAS 06475-0735 are consistent with the optically thick disc surrounded by a flared thin region, as proposed for their Group I. On the extreme is IRAS 07394-1953, showing the highest $f_c$ of the sample, which precedes the two groups of Meeus et al."
 in the sequence. since ib. is not an isolated LLXeBoe.," in the sequence, since it is not an isolated HAeBe."
 Llowever. there are no clear dilferences among the derived model parameters or these stars that were chosen to represent dilferent groups. or SED shapes.," However, there are no clear differences among the derived model parameters for these stars that were chosen to represent different groups, or SED shapes."
 Based only on the SED fitting it is dillicult to distinguish the objects in an evolutionary sequence. mainly or those showing 1κο<0.7.," Based only on the SED fitting it is difficult to distinguish the objects in an evolutionary sequence, mainly for those showing $-1 < \beta_1 < 0.7$."
 Indeed. the comparison of the stellar parameters on he lLH-I. diagram shows no evidence of cdillerences on heir evolutionary status.," Indeed, the comparison of the stellar parameters on the H-R diagram shows no evidence of differences on their evolutionary status."
 In Figure. 1 the bolometric uminosities ancl ellective temperatures of the four stars are compared with the isochrones provided by Siessοἱal. (2000)., In Figure \ref{fighr} the bolometric luminosities and effective temperatures of the four stars are compared with the isochrones provided by \cite{Siess00}.
.. The Zero Age Main Sequence (ZXMS). and isochrones for 10. 5 and 1 Mywr are clisplaved.," The Zero Age Main Sequence (ZAMS) and isochrones for 10, 5 and 1 Myr are displayed."
 Even considering the well-determined age of LID 141569. in the literature. it is not possible to infer the ages of the other objects in Figure LL due to the imprecision on distance or temperature determination. as indicated by the illustrative error bars.," Even considering the well-determined age of HD 141569 in the literature, it is not possible to infer the ages of the other objects in Figure \ref{fighr} due to the imprecision on distance or temperature determination, as indicated by the illustrative error bars."
" In spite of the age uncertainties. ib can be noted that high values of 2, correspond to higher positions above the ZAMS. but a large number of objects is needed to verily this trend."," In spite of the age uncertainties, it can be noted that high values of $\beta_1$ correspond to higher positions above the ZAMS, but a large number of objects is needed to verify this trend."
" In fact. the only hypothesis for the studied: sample is the suggestion regarding the extremities of a possible sequence,"," In fact, the only hypothesis for the studied sample is the suggestion regarding the extremities of a possible sequence."
 3)«LO objects showing very little circumstellar emission (f.<— 10%). as the example. of LID 141569. are the best. candidates to have more evolved. clises. while IUS 07394-1953 probably is the opposite case.," $\beta_1 < -1.0$ objects showing very little circumstellar emission $f_c < 10\%$ ), as the example of HD 141569, are the best candidates to have more evolved discs, while IRAS 07394-1953 probably is the opposite case."
 Any other inferences about the evolutionary status of objects require a demonstration for a large sample that we intend to develop in a further work., Any other inferences about the evolutionary status of objects require a demonstration for a large sample that we intend to develop in a further work.
 For LD 141569 the SED could not be well fitted with the disc raclius measured from NICAIOS image. indicating that the [lared. disc model fails to give the best fitting mainly in the mid-infrared. region around ~25 yam. This is not due to a failure of the GA method to find the real best solution.," For HD 141569 the SED could not be well fitted with the disc radius measured from NICMOS image, indicating that the flared disc model fails to give the best fitting mainly in the mid-infrared region around $\sim 25$ $\mu$ m. This is not due to a failure of the GA method to find the real best solution."
 Lt is verified that a different parameters set. can provide an adequate SED fitting. but resulting in clise raclius Ry=18 AU.," It is verified that a different parameters set can provide an adequate SED fitting, but resulting in disc radius $R_D = 13$ AU."
 While this seems to contradict the near-infrared observations. which revealed a cise size larger than 300 AU. it must to be noted that only sub-jum sized. grains are seen in scattered light at large distances.," While this seems to contradict the near-infrared observations, which revealed a disc size larger than 300 AU, it must to be noted that only $\mu$ m sized grains are seen in scattered light at large distances."
 The outer parts of the disc do not comprise the bulls of the dise mass. which is mostly concentrated. closer to the star.," The outer parts of the disc do not comprise the bulk of the disc mass, which is mostly concentrated closer to the star."
 Probably. the evolved disc (or thin cise gcometrv) of LID 141569 is better explained by a fat structure instead.," Probably, the evolved disc (or thin disc geometry) of HD 141569 is better explained by a flat structure instead."
 Our previous results obtained using the flat disc mocel are consistent with the observed size (see Appendix C)., Our previous results obtained using the flat disc model are consistent with the observed size (see Appendix \ref{ttapp}) ).
 llowever. that. simple model also fails to reproduce. the complicated structure of the inner disc of HD. 141569 that seems not to be only related to dillerences on circumstellar ecometry or degeneracy of the models.," However, that simple model also fails to reproduce the complicated structure of the inner disc of HD 141569 that seems not to be only related to differences on circumstellar geometry or degeneracy of the models."
 We trace below some information about grain sizes and chemical composition. probably related to evolutionary aspects.," We trace below some information about grain sizes and chemical composition, probably related to evolutionary aspects."
 The presence of large grains (Doccalettietal.2003) and Als (Sylvesteretal.1996:Weinberger2004:Geersal.2006). have been reported for HD. 141569. while silicate emission is not present. (Weinberecretal.2000).," The presence of large grains \citep{Boccaletti03} and PAHs \citep{Sylvester96, Weinberger04, Geers06} have been reported for HD 141569, while silicate emission is not present \citep{Weinberger00}."
 This can » verified in the spectral features presented in Figure 10.., This can be verified in the spectral features presented in Figure \ref{iso398}.
 The data show an almost Hat SED near 10 pam. while a ow intensity PALL emission is found at 8.6 yan. The absence of the LO yam feature has been discussed » AMeeusetal.(2002) as a possible lack of small silicate erains. or a temperature effect that leads to a lack of small. hot silicates.," The data show an almost flat SED near 10 $\mu$ m, while a low intensity PAH emission is found at 8.6 $\mu$ m. The absence of the 10 $\mu$ m feature has been discussed by \cite{Meeus02} as a possible lack of small silicate grains, or a temperature effect that leads to a lack of small, hot silicates."
 According DDNOL. the shadowing caused by the rim can suppress the strength of the silicate emission.," According DDN01, the shadowing caused by the rim can suppress the strength of the silicate emission."
 Hernándezetal.(2007) suggest that grain growth and/or settling. and transitions objects. which are stars with inner gaps in their discs (SiciliazXeguilaretal.2006).. are à possible combination to explain. evolved. discs.," \cite{Hernandez07} suggest that grain growth and/or settling, and transitions objects, which are stars with inner gaps in their discs \citep{Sicilia06}, are a possible combination to explain evolved discs."
 The correlation between the mid-infrared and the sub millimetre slope. found by Ackeetal.(2004). indicates. that as the grains grow the disc structure evolves from. Hlared. to ecometrically Hat (Dullemond2002).. which seems to be the case of LID 141569.," The correlation between the mid-infrared and the sub millimetre slope, found by \cite{Acke04}, indicates that as the grains grow the disc structure evolves from flared to geometrically flat \citep{Dullemond02}, which seems to be the case of HD 141569."
 Another implication on the radial structure of this disc can be also related to dynamical processes., Another implication on the radial structure of this disc can be also related to dynamical processes.
 Scattered light images of the optically thin dust disc were used by Augercauetal.(1999). and. Augereau&Papaloizou(2004) in the investigation of the properties and the dynamics ofthe LED 141569 triple svstem., Scattered light images of the optically thin dust disc were used by \cite{Augereau99} and \cite{Augereau04} in the investigation of the properties and the dynamics ofthe HD 141569 triple system.
 Their suggestions on different cdga populations required to fit the full SED. as well as the processes of disctruncation. indicate a scenario too complex to be reproduced by the two models that we have studied.," Their suggestions on different dust populations required to fit the full SED, as well as the processes of disctruncation, indicate a scenario too complex to be reproduced by the two models that we have studied."
sources within the field of view. the entire field of view must be imaged.,"sources within the field of view, the entire field of view must be imaged."
 Production of the images was conducted in a consistent manner [rom epoch to epoch., Production of the images was conducted in a consistent manner from epoch to epoch.
 Calibration of the flux densitv was by reference either to or236., Calibration of the flux density was by reference either to or.
. Initial calibration ol the visibility phases was obtained by observations of a nearby VLA or GAIRT calibrator. ivpicallvJ1714—252.," Initial calibration of the visibility phases was obtained by observations of a nearby VLA or GMRT calibrator, typically."
. At 330 MIIz. radio frequency interlerence (RFI) can be a substantial problem. ancl. if not excised from (he visibility data. it would limit the dynamic range of the final image.," At 330 MHz, radio frequency interference (RFI) can be a substantial problem, and, if not excised from the visibility data, it would limit the dynamic range of the final image."
 We examined the visibility data for and excised it., We examined the visibility data for and excised it.
 At 330 MIIz. neither the VLA nor the GMBRT can be assumed to be coplanar: in order {ο image the entire field of view. we used a polvhedral imaging algorithm to compensate for the non-coplanaritv of the arravs (Cornwell&Perley1992).," At 330 MHz, neither the VLA nor the GMRT can be assumed to be coplanar; in order to image the entire field of view, we used a polyhedral imaging algorithm to compensate for the non-coplanarity of the arrays \citep{cp92}."
. In order (ο approach thermal noise limits in the images. several iterations of imaging. deconvolution Niing). and were used.," In order to approach thermal noise limits in the images, several iterations of imaging, deconvolution ing), and self-calibration were used."
 In order {ο search for bursts [romJ1745—3009.. the components of all other sources in the field were subtracted [rom the «or data. ancl the residual data were then imaged in 10 min.," In order to search for bursts from, the components of all other sources in the field were subtracted from the $u$ $v$ data, and the residual data were then imaged in 10 min."
 subsets., subsets.
 Noise levels of the 10 iimages range from approximatelv 10 flor the GMRT and 20 {for the most extended VLA configurations (A and D) to approximately 250 {for the more compact configurations (C and D). which have both a lower angular resolution and are more susceptible to and siclelobe confusion.," Noise levels of the 10 images range from approximately 10 for the GMRT and 20 for the most extended VLA configurations (A and B) to approximately 250 for the more compact configurations (C and D), which have both a lower angular resolution and are more susceptible to and sidelobe confusion."
 If a burst was detected. the residual data then were imaged with a higher time resolution (from 5 (o 30J) s) in order (ο search [or siructure within the burst.," If a burst was detected, the residual data then were imaged with a higher time resolution (from 5 to 30 s) in order to search for structure within the burst."
 All but four of the VLA observations listed in Table 1. are pointed in the direction ofA*.. approximately nnorth ofJ1745—3009.. (," All but four of the VLA observations listed in Table \ref{tab:log} are pointed in the direction of, approximately north of. ("
Coincidentally. the discovery observations were pointed nearly directly at the source.),"Coincidentally, the discovery observations were pointed nearly directly at the source.)"
 The GAIRT observations are pointed approximately wwest., The GMRT observations are pointed approximately west.
 The primary beam attenuation of the VLA and the GAIRT reduces the apparent Πας density of the source by a [actor of approximately 2 ancl 1.5. respectively.," The primary beam attenuation of the VLA and the GMRT reduces the apparent flux density of the source by a factor of approximately 2 and 1.5, respectively."
 While signilicant. this level of primary. beam attenuation would not be sufficient to prevent the recovery οἱ the source. provided that (he amplitude of (he bursts is approximately 1 Jv.," While significant, this level of primary beam attenuation would not be sufficient to prevent the recovery of the source, provided that the amplitude of the bursts is approximately 1 Jy."
 ILowever. if the bursts have a range of amplitudes. significantly weaker bursts (S150 mJv) could have gone undetected in the vast majority of our observations.," However, if the bursts have a range of amplitudes, significantly weaker bursts $\lesssim 150$ mJy) could have gone undetected in the vast majority of our observations."
source GNX 13|LI. based on the estimated distance and Ny to cach source.,"source GX 13+1, based on the estimated distance and ${\rm N_H}$ to each source."
 There is a rather large range. from as bright as 3.8 for GX 131110 possidy as faint as 2.9 for GX 1712 if the observed. IX magnitude in quiescence is as faint as ~ISS (Callanan et al.," There is a rather large range, from as bright as $-3.8$ for GX 13+1 to possibly as faint as $+2.9$ for GX 17+2 if the observed K magnitude in quiescence is as faint as $\sim 18.5$ (Callanan et al."
 1999)., 1999).
 Uncertainties in distance estimates and reddening are likely to be significant at a evel of about £1 magnitude. and so cannot account for the ooad range.," Uncertainties in distance estimates and reddening are likely to be significant at a level of about $\pm 1$ magnitude, and so cannot account for the broad range."
 Several dillerent. components may. contribute significantly to the emission in the near.infrared: as a uide o their significance (see below) we have also listed. binary orbial periods and. where available. the spectral types of he mass donors in ‘Table 4.," Several different components may contribute significantly to the emission in the near–infrared; as a guide to their significance (see below) we have also listed binary orbital periods and, where available, the spectral types of the mass donors in Table \ref{zsources}."
 Vhermal emission. will be LOCuced both by the stellar companion and the accretion disc (for a discussion. of their. relative contributions see also Bandyopadhyay et a., Thermal emission will be produced both by the stellar companion and the accretion disc (for a discussion of their relative contributions see also Bandyopadhyay et al.
 1997: 1090)., 1997; 1999).
 We may expect the accreion-dise contribution to depend on the size of the disc (van Paradijs AleClintock 1994). which in urn should 20 0 unction of the orbital period of the system.," We may expect the accretion-disc contribution to depend on the size of the disc (van Paradijs McClintock 1994), which in turn should be a function of the orbital period of the system."
 We note hat for the three svstems with some attempt at spectral classification of the mass donor there is a good agreement oetween the absolute Ix band magnitudes derived and those expected for the companion spectral class., We note that for the three systems with some attempt at spectral classification of the mass donor there is a good agreement between the absolute K band magnitudes derived and those expected for the companion spectral class.
 Εμ implies hat GX 51 shoud contain a relatively bright mass donor., This implies that GX 5–1 should contain a relatively bright mass donor.
 Luminosity class HE was found for the companion star in Sco X Land Cvg X2 (see Table. 4.. and references therein)," Luminosity class III was found for the companion star in Sco X–1 and Cyg X–2 (see Table \ref{zsources}, and references therein)."
 Following the conjecture made w Hasinger van der Ixlis (1989) that all Z sources have evolved: companions. we assume also luminosity class LLL or GN 51.," Following the conjecture made by Hasinger van der Klis (1989) that all Z sources have evolved companions, we assume also luminosity class III for GX 5–1."
 The companion star in GX 51 is hen most likely of spectral type Ix. There may also be an additional contribution from infrared: svnchrotron. emission. as found in the black hole system GRS 1915105 (Fender Pooley 1998 anc references herein).," The companion star in GX 5–1 is then most likely of spectral type K. There may also be an additional contribution from infrared synchrotron emission, as found in the black hole system GRS 1915+105 (Fender Pooley 1998 and references therein)."
 Wat all. this should onlv occur when tie source is radiobright.," If at all, this should only occur when the source is radio–bright."
 Phe Z sources are brightest at racio wavelenghs when thev are observed on the Horizontal Branch (IB) in the X-ray coour-colour diagram. (DPenninx et al., The Z sources are brightest at radio wavelengths when they are observed on the Horizontal Branch (HB) in the X-ray colour-colour diagram (Penninx et al.
 1988: IHjellming Llan 1995)., 1988; Hjellming Han 1995).
 Radio Waring in Z sources typically has ampliudes of a [ow mJv (οαπο Lan 1995 and references. therein): if the synchrotron emission has a flat spectrum to the near-infrared we might observe a (reddened) contribution of <1 niJy at times., Radio flaring in Z sources typically has amplitudes of a few mJy (Hjellming Han 1995 and references therein); if the synchrotron emission has a flat spectrum to the near-infrared we might observe a (reddened) contribution of $\ga 1$ mJy at times.
 For 6X 5.1 this could cause up to 1 magnitude variability., For GX 5–1 this could cause up to 1 magnitude variability.
 If star 513 is not the counterpart of GX 51. the counterpart must have been 22.5 magnitudes fainter in the Ix band at the time of our observations.," If star 513 is not the counterpart of GX 5–1, the counterpart must have been $\ga$ 2.5 magnitudes fainter in the K band at the time of our observations."
 Future spectroscopic observations and/or the detection of variability should confirm star 513 as the counterpart., Future spectroscopic observations and/or the detection of variability should confirm star 513 as the counterpart.
 ‘To conclude. we have most likely identified the Lh Counterxut of the bright Z-twpe X-ray source GX 5.1 based upon positional coincidence with the radio counterpart. an identification which is supported by marginal evidence for excess Bre emission.," To conclude, we have most likely identified the IR counterpart of the bright Z-type X-ray source GX 5–1 based upon positional coincidence with the radio counterpart, an identification which is supported by marginal evidence for excess $\gamma$ emission."
 We have discussed the possible origins of LR emission in this system and in the other Z sources (and GX 13]1). and suggest that GX 51 may contain a WLLL mass conor.," We have discussed the possible origins of IR emission in this system and in the other Z sources (and GX 13+1), and suggest that GX 5–1 may contain a KIII mass donor."
 The UBER is operated by the Joint Astronomy Centre on behalfol the U.Ix. Particle Physies and Astronomy Research Council., The UKIRT is operated by the Joint Astronomy Centre on behalf of the U.K. Particle Physics and Astronomy Research Council.
 Phe data reported here were obtained as part of the UBIRT Service Programme., The data reported here were obtained as part of the UKIRT Service Programme.
 We would like to thank Sandy Leggett for the Alay 23 observations and helpful comments in reducing the observations. John Davies for the October observations. Paul Vreeswijk for help with the reduction of the images. Tim Navlor for providing us with electronic versions of the images presented. in Navlor ct al. (," We would like to thank Sandy Leggett for the May 23 observations and helpful comments in reducing the observations, John Davies for the October observations, Paul Vreeswijk for help with the reduction of the images, Tim Naylor for providing us with electronic versions of the images presented in Naylor et al. ("
1991) facilitating comparisons. and the referee. Phil Charles for his comments which improved the paper.,"1991) facilitating comparisons, and the referee, Phil Charles for his comments which improved the paper."
similar SEIt.,similar SFR.
 This indicates that. once the heating epoch is fixed. the degree of SER. suppression depends only on the amount of heating energy. while being largely. independen on its clistribution as a function of the local gas density.," This indicates that, once the heating epoch is fixed, the degree of SFR suppression depends only on the amount of heating energy, while being largely independent on its distribution as a function of the local gas density."
 Llowever. heating at z=9 with a comparable amount of energy does not allow gas to reach high densities within DN halos and to cool before z21.," However, heating at $z=9$ with a comparable amount of energy does not allow gas to reach high densities within DM halos and to cool before $z\simeq 1$."
 Once cooling takes place. 1 converts less than 10 per cent ofthe gas into stars. within à short episode of star formation.," Once cooling takes place, it converts less than $10$ per cent of the gas into stars, within a short episode of star formation."
 Ehe resulting SER peaks a very low redshift. 20.3. which is highly discrepant. with observational determinations of the SER. history in clusters (ο... Ixodama Bower 2001).," The resulting SFR peaks at very low redshift, $z\simeq
0.3$, which is highly discrepant with observational determinations of the SFR history in clusters (e.g., Kodama Bower 2001)."
 Of course. this is not the only feature which rules out the picture of a strong heating occurring at such a high redshift.," Of course, this is not the only feature which rules out the picture of a strong heating occurring at such a high redshift."
 For instance. at z0. stars are all concentrated in one single object located at the center. a cDlike galaxy. while no other galaxysized DM halos contains significantD amounts of collapsed 5gas.," For instance, at $z=0$, stars are all concentrated in one single object located at the center, a cD–like galaxy, while no other galaxy–sized DM halos contains significant amounts of collapsed gas."
 As for the SN heating (right panel. of Fig. 49).," As for the SN heating (right panel of Fig. \ref{fi:sfr}) ),"
 suppressing the star fraction below the 10 per cent level requires a high. probably unrealistic value. for ro.," suppressing the star fraction below the 10 per cent level requires a high, probably unrealistic value for $\eta_0$."
" Also. taking go=1.510TM* generates an implausible SER history. resembling that found for the runs based on setting the entropy floor at 2,=9:IR the large amount of extra heating at high redshift prevents the occurrence of star formation down to z1.5."," Also, taking $\eta_0=1.5\times 10^{-2}M_\odot^{-1}$ generates an implausible SFR history, resembling that found for the runs based on setting the entropy floor at $z_h=9$: the large amount of extra heating at high redshift prevents the occurrence of star formation down to $z\simeq
1.5$."
 Taking go in the range 10. predicts more realistic ος. but it is not able to suppress foot below the 220 per cent level.," Taking $\eta_0$ in the range $\times 10^{-3}$ predicts more realistic SFRs, but it is not able to suppress $f_{\rm cold}$ below the $\simeq 20$ per cent level."
 A general conclusion of our analvsis is that. heating schemes producing plausible SER. histories are not efficient in suppressing the fraction of cold gas below the 20 and 25 per cent values at the cluster and group scales. respectively.," A general conclusion of our analysis is that heating schemes producing plausible SFR histories are not efficient in suppressing the fraction of cold gas below the 20 and 25 per cent values at the cluster and group scales, respectively."
 Vice-versa. à more efficient. suppression is obtained. by preventing gas to cool at high. recdshi at the expense of delaving star formation to unreasonably low redshilts.," Vice-versa, a more efficient suppression is obtained by preventing gas to cool at high redshift, at the expense of delaying star formation to unreasonably low redshifts."
 Aleasurements of the excess entropy in central regions of »o»or clusters and &roups are considered. to provide clirec evidence for the lack of self.similarity of the ICM properties (c.g.. Ponman et al.," Measurements of the excess entropy in central regions of poor clusters and groups are considered to provide direct evidence for the lack of self–similarity of the ICM properties (e.g., Ponman et al."
 1999: Finoguenov ct al., 1999; Finoguenov et al.
 2002a)., 2002a).
 In a separate paper (Finoguenov ct al., In a separate paper (Finoguenov et al.
 2002b). a selfconsisten comparison is realized. between the entropy properties. of he simulations with impulsive heating. that we presen rere. and the observational data for groups and. clusters o» Linoguenov et al. (," 2002b), a self–consistent comparison is realized between the entropy properties of the simulations with impulsive heating, that we present here, and the observational data for groups and clusters by Finoguenov et al. ("
2002a).,2002a).
 The main result of this comparison is that. although cooling ancl star formation end to somewhat increase entropy in central cluster regions. hey still fall short in producing the entropy excess which is observedbserved at the group scale.," The main result of this comparison is that, although cooling and star formation tend to somewhat increase entropy in central cluster regions, they still fall short in producing the entropy excess which is observed at the group scale."
" While prelpreheating eatat ο=3j is shown to increase the entropy to the observed values. runs with 2,=9 are characterized by a low entropy. level in central regions of clusters ancl groups."," While preheating at $z_h=3$ is shown to increase the entropy to the observed values, runs with $z_h=9$ are characterized by a low entropy level in central regions of clusters and groups."
 Insteacl of attempting any further comparison with observations. we want to discuss here the dynamical reasons for such a behavior.," Instead of attempting any further comparison with observations, we want to discuss here the dynamical reasons for such a behavior."
" To this end. we show in Figure 6 the cllect of cooling and nongravitational heating on the entropy profiles for our whole set of ""Virgo simulations."," To this end, we show in Figure \ref{fi:sprof} the effect of cooling and non–gravitational heating on the entropy profiles for our whole set of “Virgo” simulations."
 As expected. when cooling anc SE are included. low entropy eas is selectively removed in central cluster regions. thus inducing a Ilattening of the profile.," As expected, when cooling and SF are included, low entropy gas is selectively removed in central cluster regions, thus inducing a flattening of the profile."
 Phis is explicitly shown in Figure. 7: while the run including only gravitational wating has a population of highdensity low entropy gas xwticles. such particles are removed from the dilfuse phase once cooling and star formation are introduced.," This is explicitly shown in Figure \ref{fi:srho}: while the run including only gravitational heating has a population of high–density low entropy gas particles, such particles are removed from the diffuse phase once cooling and star formation are introduced."
 Fhis result is consistent with the expectation from analvtical arguments xwed on the comparison between cooling timescale and vpical cluster age (e.g... Voit et al.," This result is consistent with the expectation from analytical arguments based on the comparison between cooling time–scale and typical cluster age (e.g., Voit et al."
 2002. Wu Xue 2002).," 2002, Wu Xue 2002)."
 The inclusion of extra heating has a nontrivial effect. on he efficiency of cooling in removing particles from the lower, The inclusion of extra heating has a non–trivial effect on the efficiency of cooling in removing particles from the lower
as already pointed out in the introduction. derived the GSMF by using a different procedure. te. by combining the UV luminosity function with an average M./L ratio.,"as already pointed out in the introduction, derived the GSMF by using a different procedure, i.e. by combining the UV luminosity function with an average $M_*/L$ ratio."
 The systematic uncertainties caused by the various assumptions involved in spectral energy distribution modelling were shown to dominate the overall error budget affecting the GSMF (seeMarchesinietal.2009.foradetailed analysis).., The systematic uncertainties caused by the various assumptions involved in spectral energy distribution modelling were shown to dominate the overall error budget affecting the GSMF \citep[see][for a detailed analysis]{marchesini09}.
 In this regard. a significant role is played by the choice of the stellar templates used to estimate the stellar mass.," In this regard, a significant role is played by the choice of the stellar templates used to estimate the stellar mass."
 Stellar masses obtained using the CBO7 stellar library. which includes an improved TP-AGB stars treatment. are on average 0.12 dex lower than those inferred using the BCO3 templates. with a scatter as large as 0.17 dex.," Stellar masses obtained using the CB07 stellar library, which includes an improved TP-AGB stars treatment, are on average 0.12 dex lower than those inferred using the BC03 templates, with a scatter as large as 0.17 dex."
 We plot in Fig., We plot in Fig.
 4 their ratio as a function of the stellar mass adopted as a reference in this work (ήμουν) in different redshift bins., \ref{fig:cfrmass} their ratio as a function of the stellar mass adopted as a reference in this work $M_{BC03}$ ) in different redshift bins.
 The lack of a clear trend of Mpcoi/Mcpos with stellar mass or redshift translates into a lack of a rigid offset between the GSMFs computed with the two libraries. although the CBO7 points are on average at lower stellar masses than BCO3.," The lack of a clear trend of $M_{BC03}/M_{CB07}$ with stellar mass or redshift translates into a lack of a rigid offset between the GSMFs computed with the two libraries, although the CB07 points are on average at lower stellar masses than BC03."
 We compare in Fig., We compare in Fig.
 5 the GSMFs obtained with the BCO3 templates (black solid curves/solid circles) and the CBO7 ones (red dotted curves/open boxes)., \ref{fig:bc03cb07} the GSMFs obtained with the BC03 templates (black solid curves/solid circles) and the CB07 ones (red dotted curves/open boxes).
 For the sake of simplicity. we decided to report the four most representative bins.," For the sake of simplicity, we decided to report the four most representative bins."
 The results for the 1.0—1.4 and 1.4—1.8 redshift bins are very similar to the 0.6—1.0 and 1.8—2.5 ones. respectively.," The results for the $1.0-1.4$ and $1.4-1.8$ redshift bins are very similar to the $0.6-1.0$ and $1.8-2.5$ ones, respectively."
" We also show the 1/V,,. points of Marchesinietal.(2009) (their set 8) and Caputietal.(2011).. both obtained by adopting the CBO7 templates."," We also show the $1/V_{max}$ points of \cite{marchesini09} (their set 8) and \cite{caputi11}, both obtained by adopting the CB07 templates."
 The results of Marchesinietal.(2009) agree with our CBO7-based GSMF in all except the 1.8—2.5 redshift interval. likely because of imperfect redshift overlap between the two analysis.," The results of \cite{marchesini09} agree with our CB07-based GSMF in all except the $1.8-2.5$ redshift interval, likely because of imperfect redshift overlap between the two analysis."
 The points from Caputiet are in broad agreement with ours at the bright end. while the incompleteness that the authors claim to be affected by below M.~10''M. is likely responsible for the disagreement at low stellar masses.," The points from \cite{caputi11} are in broad agreement with ours at the bright end, while the incompleteness that the authors claim to be affected by below $M_*\sim 10^{11} M_\odot$ is likely responsible for the disagreement at low stellar masses."
 The best-fit Schechter parameters of the CBO7-based GSMF are reported in Tab. 2.., The best-fit Schechter parameters of the CB07-based GSMF are reported in Tab. \ref{tab:paramcb07}.
 At z>2.5. we were forced to the M parameter to its best-fit value at z~2.15.," At $z>2.5$, we were forced to the $M^*$ parameter to its best-fit value at $z\sim 2.15$."
 If it is instead allowed to vary. the fit is unconstrained or the maximum-likelihood analysis does not converge.," If it is instead allowed to vary, the fit is unconstrained or the maximum-likelihood analysis does not converge."
 The CBO7- and BCO3-based GSMEFs differ from each other., The CB07- and BC03-based GSMFs differ from each other.
 However. we do not find any similar systematic behaviour at all redshifts.," However, we do not find any similar systematic behaviour at all redshifts."
That the high-mass end of the CBO7-based GSMF is unconstrained at >2.5. while the BCO3-based one suffers from poor statistical sampling only in the highest redshift bin (:> 3.5). is à further confirmation that the two GSMFs are not affected by a systematic shift in stellar mass.,"That the high-mass end of the CB07-based GSMF is unconstrained at $z>2.5$, while the BC03-based one suffers from poor statistical sampling only in the highest redshift bin $z>3.5$ ), is a further confirmation that the two GSMFs are not affected by a systematic shift in stellar mass."
 At— the lowest and the highest redshifts. we find the closer agreement. the normalization of the best-fit Schechter function being only slightly lower when CBO7 templates are used.," At the lowest and the highest redshifts, we find the closer agreement, the normalization of the best-fit Schechter function being only slightly lower when CB07 templates are used."
 At intermediate redshifts. we observe a more serious disagreement. resulting in different faint-end slopes and characteristic masses.," At intermediate redshifts, we observe a more serious disagreement, resulting in different faint-end slopes and characteristic masses."
 This is unsurprising. because the effect of including of the TP-AGB phase is expected to be important at intermediate ages (0.2-—2 Gyr). which predominate the 2€zx3 redshift range (Maraston2005:Henriquesetal.2011).," This is unsurprising, because the effect of including of the TP-AGB phase is expected to be important at intermediate ages $0.2 - 2$ Gyr), which predominate the $2 \lesssim z \lesssim 3$ redshift range \citep{maraston05,henriques11}."
. Although the difference between the CBO7- and the BCO3-based GSMFs do not show a systematic. trend at all redshifts. the characteristic masses seem to be on average lower when CBO7 templates are used (see Fig.," Although the difference between the CB07- and the BC03-based GSMFs do not show a systematic trend at all redshifts, the characteristic masses seem to be on average lower when CB07 templates are used (see Fig."
 6 discussed in the next section). as expected. despite the large uncertainties. in agreement with the results of Marchesinietal.," \ref{fig:contours} discussed in the next section), as expected, despite the large uncertainties, in agreement with the results of \cite{marchesini09} ."
(2009) This trend is clear at 1.4.«z2.5. where our redshift bins overlap with those of Marchesinietal.(2009)... while the lack of high quality statistics at higher redshifts prevents us from drawing any firm conclusions about the effect ofchanging the stellar templates.," This trend is clear at $1.4 < z < 2.5$, where our redshift bins overlap with those of \cite{marchesini09}, while the lack of high quality statistics at higher redshifts prevents us from drawing any firm conclusions about the effect ofchanging the stellar templates."
 For what concerns the variation in a when changing the template library. we find similar slopes from zo0.8 t0z- 1.2. while in the redshift interval 1.4—2.5 the BCO3-based GSMFs are steeper than the CBOT7-based ones by 0.2—0.3.," For what concerns the variation in $\alpha$ when changing the template library, we find similar slopes from $z\sim 0.8$ to $z\sim 1.2$, while in the redshift interval $1.4- 2.5$ the BC03-based GSMFs are steeper than the CB07-based ones by $0.2-0.3$."
 Marchesinietal.(2009) reports similar slopes when using BC03 and CBO7 stellar templates in the redshift interval 1.3-3.0. while their BCO3-based GSME is steeper than the CBO7-based one at 3.0«z4.0.," \cite{marchesini09} reports similar slopes when using BC03 and CB07 stellar templates in the redshift interval $1.3 - 3.0$, while their BC03-based GSMF is steeper than the CB07-based one at $3.0<z<4.0$."
 The intrinsic difference between the two surveys does not allow us to investigate the origin of this mismatch., The intrinsic difference between the two surveys does not allow us to investigate the origin of this mismatch.
 The main goal of this study has been to investigate the faint-end slope of the GSME. especially at the highest redshifts (z>2).," The main goal of this study has been to investigate the faint-end slope of the GSMF, especially at the highest redshifts $z>2$ )."
 From both Fig., From both Fig.
 3 and Tables | and 2.. it is evident that the low- tail steepens with redshift.," \ref{fig:mf_obs} and Tables \ref{tab:parambc03} and \ref{tab:paramcb07}, it is evident that the low-mass tail steepens with redshift."
 The results from applying the STY approach to our BCO3-based data indicate that the faint- slope steepens significantlybetween z~0.8. where we fitted vw=—1.44x0.03. and z~ 3. where the best-fit « Is equal to —1.56+ 0.16. before flattening up to z~ 4.," The results from applying the STY approach to our BC03-based data indicate that the faint-end slope steepens significantlybetween $z\sim 0.8$, where we fitted $\alpha = -1.44 \pm 0.03$, and $z\sim 3$ , where the best-fit $\alpha$ is equal to $-1.86 \pm 0.16$ , before flattening up to $z\sim 4$ ."
 First of all. we performed a simple sanity check to verify that the abundance of low mass objects at z>1.8 is reliable by plotting all objects with M.«I0'M. and 1.8«z2.5 on à BzK diagram.," First of all, we performed a simple sanity check to verify that the abundance of low mass objects at $z > 1.8$ is reliable by plotting all objects with $M_* < 10^{10}M_\odot$ and $1.8<z<2.5$ on a $BzK$ diagram."
 For galaxies at z>2.5. we adopted the analogous RJL diagram (using IRAC 3.6 as L band). which extends the former to the 2.5<z«4 redshift regime (Daddietal. 2004).. and checked stellar masses below 2-10A...," For galaxies at $z>2.5$, we adopted the analogous $RJL$ diagram (using IRAC 3.6 as $L$ band), which extends the former to the $2.5<z<4$ redshift regime \citep{daddi04}, and checked stellar masses below $ 2 \cdot 10^{10}M_\odot$."
 Approximately of the sources indeed lie in the high redshift regions of these diagrams. makingus confident of their photometric redshift estimate.," Approximately of the sources indeed lie in the high redshift regions of these diagrams, makingus confident of their photometric redshift estimate."
" As an additional check. we carefully inspected the individual photometric-redshift probability distribution curves for each source with z>1.8 and M.« 2-10!'""M... Following the criterion described in Sect. 2.. "," As an additional check, we carefully inspected the individual photometric-redshift probability distribution curves for each source with $z> 1.8$ and $M_* <2 \cdot 10^{10}M_\odot$ Following the criterion described in Sect. \ref{sec:data}, ,"
we found that of these sources have a probability of lying at z> 1.5., we found that of these sources have a probability of lying at $z>1.5$ .
 As alreadypointed outin Sect. 3.2.. ," As alreadypointed outin Sect. \ref{sec:gsmf}, ,"
the small sky area sampled by our data may be responsible for degeneracies between the faint-end slope a and the characteristic mass M when fitting a Schechter function., the small sky area sampled by our data may be responsible for degeneracies between the faint-end slope $\alpha$ and the characteristic mass $M^*$ when fitting a Schechter function.
 We therefore studiedin, We therefore studiedin
"A fact⋅ usually assumed. in. astrophysics ⊀⊀is that the main. part of the mass of a typical spiral galaxy is concentrated in. a thin ...disk (Binney""&s‘Tremaine:- (1987))).",A fact usually assumed in astrophysics is that the main part of the mass of a typical spiral galaxy is concentrated in a thin disk \cite{BT}) ).
" Accordingly.. the obtention of the eravitational rotlcntial generated by an idealized thin disk is a problem of great. astrophysical relevance and so. through the vears. cillerent approaches has been used to obtain such kind ο ""thin clisk mocdels."," Accordingly, the obtention of the gravitational potential generated by an idealized thin disk is a problem of great astrophysical relevance and so, through the years, different approaches has been used to obtain such kind of thin disk models."
 Wvse and AMavall (1942)) sucliccL thin clisks bv superposing an infinite family of elementary disks of dillerent. radii., Wyse and Mayall \citeyear{WM}) ) studied thin disks by superposing an infinite family of elementary disks of different radii.
 Drandt (1960)) and Brandt. and Alton (1962)) constructed. Lat galaxy elisks bv the flattening of a distribution of matter whose surface of equal density were similar spheroids., Brandt \citeyear{BR}) ) and Brandt and Belton \citeyear{BB}) ) constructed flat galaxy disks by the flattening of a distribution of matter whose surface of equal density were similar spheroids.
 simple potentia-density pair for a thin disk model was introduced. by Ixuzmin. (1952)) and. then reclerived by Toomre (1963: 19649) as the first member of a generalized Family. of moclels.," A simple potential-density pair for a thin disk model was introduced by Kuzmin \citeyear{KUZ}) ) and then rederived by Toomre \citeyear{T1,T2}) ) as the first member of a generalized family of models."
 ‘The Toomre models are obtainec Γον solving the Laplace equation in evlindrieal coordinates subject to appropriated boundary conditions on the disk and at infinity., The Toomre models are obtained by solving the Laplace equation in cylindrical coordinates subject to appropriated boundary conditions on the disk and at infinity.
 The Kuzmin and ‘Toone mocels of thin disks. although they have surface densities and rotation curves with remarkable properties. represent clisks of infinite extension and thus they are rather poor Uat galaxy mocels.," The Kuzmin and Toomre models of thin disks, although they have surface densities and rotation curves with remarkable properties, represent disks of infinite extension and thus they are rather poor flat galaxy models."
 Accordingly. in order to obtain more realistic mocels of [lat galaxies. is better to consider methods that permit the obtention of finite thin disk models.," Accordingly, in order to obtain more realistic models of flat galaxies, is better to consider methods that permit the obtention of finite thin disk models."
 A simple. method to obtain thve surface density. the gravitational potential and the rotation curve of thin disks of finite radius was developed by Hunter.(1963).," A simple method to obtain the surface density, the gravitational potential and the rotation curve of thin disks of finite radius was developed by \cite{HUN1}."
 The llunter method. is based in the obtention of solutions of Laplace equation D.in terms of⋅ oblate spheroidal⊀ coordinates.⊀ which are ideally suited to. the sudy of [at disks. of finite extension.," The Hunter method is based in the obtention of solutions of Laplace equation in terms of oblate spheroidal coordinates, which are ideally suited to the study of flat disks of finite extension."
 By superposition of solutions of Laplace equation. expressions for the surface. density of the disks. the gravitational potential and its rotational velocity can be obtained as series of elementary. functions.," By superposition of solutions of Laplace equation, expressions for the surface density of the disks, the gravitational potential and its rotational velocity can be obtained as series of elementary functions."
 The simplest example of a thin clisk obtained by means of the Hunter method. is the well known lxalnajs csk (Ixalnajs (1972))). which can also be obtained by Lattening a uniformly rotating spheroid (Mvse&Mayvall(1942):Brandt(1960):Brandt&Belton (1962))).," The simplest example of a thin disk obtained by means of the Hunter method is the well known Kalnajs disk \cite{KAL}) ), which can also be obtained by flattening a uniformly rotating spheroid \cite{WM,BR,BB}) )."
 The Ixalnajs disk have a well behavecl surface density and. represents a uniformly rotating disk. so that its circular velocity is proportional to he raclius. and its stability properties have been extensively studied: (see. for instance. Hunter. (1963: 1965)). Ixalnajs(1972) and Walnajs&Athanassoula-Georgala (1974))).," The Kalnajs disk have a well behaved surface density and represents a uniformly rotating disk, so that its circular velocity is proportional to the radius, and its stability properties have been extensively studied (see, for instance, Hunter \citeyear{HUN1,HUN2}) ), \cite{KAL} and \cite{K-A}) )."
 In this paper we use the Hunter method. in order to obtain an infinite family of thin disks of finite racdus., In this paper we use the Hunter method in order to obtain an infinite family of thin disks of finite radius.
 We xwticularize the Llunter general model by considering a zumilv of thin clisks with a well behaved surface mass density., We particularize the Hunter general model by considering a family of thin disks with a well behaved surface mass density.
 We will require that the surface density be a monotonically decreasing function of the radius. with à maximum at the center of the disk and vanishing at the edge. in such a wav that the mass distribution of the higher members of the family be more concentrated at the center.," We will require that the surface density be a monotonically decreasing function of the radius, with a maximum at the center of the disk and vanishing at the edge, in such a way that the mass distribution of the higher members of the family be more concentrated at the center."
 The paper is organized as follows., The paper is organized as follows.
 In Sec., In Sec.
 2 we present a summary of the Llunter method: used to obtain the thin disk mocels of finite radius and also we obtain the general, 2 we present a summary of the Hunter method used to obtain the thin disk models of finite radius and also we obtain the general
It the class contaius 200 students. this cau literally be a several hour activity.,"If the class contains 200 students, this can literally be a several hour activity."
 Iustead. a short seript can be written to login into the website. grab a file of student grades on the local machine aud POST them to the website by studeut ID (although your local network services might be curious on how you entered 200 grades in 3.5 microseconds).," Instead, a short script can be written to login into the website, grab a file of student grades on the local machine and POST them to the website by student ID (although your local network services might be curious on how you entered 200 grades in 3.5 microseconds)."
 Again. this technique is open to abuse aud a responsible user would limit the number of interactions with a website. aud their [requency. (e.g. placing the conmunaud between page requests).," Again, this technique is open to abuse and a responsible user would limit the number of interactions with a website, and their frequency (e.g., placing the command between page requests)."
 Websites have become iucreasiugly complex iu recent vears. often with complicated cookies that the and modules fail to handle.," Websites have become increasingly complex in recent years, often with complicated cookies that the and modules fail to handle."
 However. every website must interact. with a browser. so nothing cau be eucoded or hiddeu that can't be parsed by whichever browser the user selects.," However, every website must interact with a browser, so nothing can be encoded or hidden that can't be parsed by whichever browser the user selects."
 the Ultimately.best script is one that beliaves like a browser and cau be trained to proceed to the pages internalof interest (ie. clicking the buttous).," Ultimately, the best script is one that behaves like a browser and can be trained to proceed to the internal pages of interest (i.e. clicking the buttons)."
 This is the job of the 1inodule (wwwsearch.sourcelorge.net/mechanize)., This is the job of the module (wwwsearch.sourceforge.net/mechanize).
 There are uumerous examplesat the website. but the following is a simple use in a script: Note that the website may reject User-Agents that are not a known browsers.," There are numerous examples at the website, but the following is a simple use in a script: Note that the website may reject User-Agent's that are not a known browsers."
 This script must be trained. in the sense that the user probably ueeds to manually follow tle website patlis first. then copy those paths into the seript.," This script must be trained, in the sense that the user probably needs to manually follow the website paths first, then copy those paths into the script."
 Aud more detective work is probably required on the FORM variables and their usage., And more detective work is probably required on the FORM variables and their usage.
 The ultimate goal for a script that uses mechauize is to parse and understaud what a webpage meaus. aud use that information to make decisions.," The ultimate goal for a script that uses mechanize is to parse and understand what a webpage means, and use that information to make decisions."
 This would form the front end of a thiukiug or knowledge system. one that harvests information at a higher level than just reducing the data from tabular form.," This would form the front end of a thinking or knowledge system, one that harvests information at a higher level than just reducing the data from tabular form."
 This will be the focus of our future work., This will be the focus of our future work.
region around the phase centre.,region around the phase centre.
 In this general. non-coplanar case. the baseline vectors u are considered to be the separatio vectors between the pairs of antennas projected onto a plan[2 normal to the phase centre.," In this general, non-coplanar case, the baseline vectors $\mathbf{u}$ are considered to be the separation vectors between the pairs of antennas projected onto a plane normal to the phase centre."
 The direction cosines / and a ar[7 also in this case taken in à basis in which (7.7i)=(0.0) lies 1 the direction of the phase centre. rather than the zenith.," The direction cosines $l$ and $m$ are also in this case taken in a basis in which $(l,m)=(0,0)$ lies in the direction of the phase centre, rather than the zenith."
 If the sampling function (which ts. to first approximation. à sum of delta functions) is denoted by S then the output of the interferometer is VS.," If the sampling function (which is, to first approximation, a sum of delta functions) is denoted by $S$ then the output of the interferometer is $VS$ ."
" The Fourier inversion &!(VS) gives DzIxsB where D. known as the dirty image. is a convolution of the true sky image / with the ""dirty beam’ B=X(S)."," The Fourier inversion $\mathfrak{F}^{-1}(VS)$ gives $D=I \ast B$ where $D$, known as the dirty image, is a convolution of the true sky image $I$ with the `dirty beam' $B=\mathfrak{F}^{-1}(S)$."
 B plays the same role for an interferometer as the point-spread function of a traditional telescope., $B$ plays the same role for an interferometer as the point-spread function of a traditional telescope.
 In general. the larger the number of samples of V. the more compact the distribution of flux density in. B. and therefore the closer the correspondence between D and the true sky 7: the sensitivity of the observation will also be increased.," In general, the larger the number of samples of $V$, the more compact the distribution of flux density in $B$, and therefore the closer the correspondence between $D$ and the true sky $I$ ; the sensitivity of the observation will also be increased."
 An interferometer containing N antennas will generate N(N—1)/2 independent samples of V at each observation., An interferometer containing $N$ antennas will generate $N(N - 1)/2$ independent samples of $V$ at each observation.
 The number of samples can be further expanded via the technique of Earth-rotation synthesis., The number of samples can be further expanded via the technique of Earth-rotation synthesis.
 In this. the rotation of the Earth over the course of a day is used to generate a sequence of different baselines u for each antenna pair.," In this, the rotation of the Earth over the course of a day is used to generate a sequence of different baselines $\mathbf{u}$ for each antenna pair."
 Another way to increase the number of samples is to observe at several frequencies. a technique known as frequency synthesis.," Another way to increase the number of samples is to observe at several frequencies, a technique known as frequency synthesis."
 Because baselines are expressed in wavelengths. the fixed spatial separation between à pair of antennas generates baseline vectors u of different length at different frequencies.," Because baselines are expressed in wavelengths, the fixed spatial separation between a pair of antennas generates baseline vectors $\mathbf{u}$ of different length at different frequencies."
 Implicit in these two synthesis techniques is the assumption that the sky brightness will be constant over time in the first case and over frequency in the second., Implicit in these two synthesis techniques is the assumption that the sky brightness will be constant over time in the first case and over frequency in the second.
 A more detailed description of the fundamentals of interferometry can be found in many sources. for example Thompson. Moran Swenson (2001)) or Taylor. Carilli Perley (1999)).," A more detailed description of the fundamentals of interferometry can be found in many sources, for example Thompson, Moran Swenson \cite{thompson}) ) or Taylor, Carilli Perley \cite{syn_im_2}) )."
 The CLEAN algorithm was invented by Hóggbom (1974)) and has been further elaborated by Clark (1980)) and Cotton and Schwab (Schwab 1984)) among others., The CLEAN algorithm was invented by Höggbom \cite{hoegbom}) ) and has been further elaborated by Clark \cite{clark}) ) and Cotton and Schwab (Schwab \cite{schwab}) ) among others.
 The essence of the algorithm is to perform many iterations of a process in which a small amount of the dirty beam is subtracted. centred at the highest remaining. point in the dirty image. ideally until nothing remains in this image but noise.," The essence of the algorithm is to perform many iterations of a process in which a small amount of the dirty beam is subtracted, centred at the highest remaining point in the dirty image, ideally until nothing remains in this image but noise."
" The positions and amounts subtracted are recorded as ""clean components! and used afterwards to reconstruct an approximation to the true sky image 7.", The positions and amounts subtracted are recorded as `clean components' and used afterwards to reconstruct an approximation to the true sky image $I$.
 The gain factor by which the dirty beam i$ multiplied. and the number of iterations to perform. are parameters which are chosen ahead of time by the user.," The gain factor by which the dirty beam is multiplied, and the number of iterations to perform, are parameters which are chosen ahead of time by the user."
 CLEAN was originally presented as an empirical algorithm which appeared to produce results. although it required some experience to judge how best to apply it.," CLEAN was originally presented as an empirical algorithm which appeared to produce results, although it required some experience to judge how best to apply it."
 It is known not to perform well when applied to extended objects (Cornwell 1983)). although à modified algorithm has been shown to vield improvements here (Steer. Dewdney Itoh 1984)).," It is known not to perform well when applied to extended objects (Cornwell \cite{cornwell_1983}) ), although a modified algorithm has been shown to yield improvements here (Steer, Dewdney Itoh \cite{steer}) )."
 CLEAN ts still in wide use as a practical method of removing sampling artifacts from interferometry images., CLEAN is still in wide use as a practical method of removing sampling artifacts from interferometry images.
" The reasons for CLEAN's success were unclear for some time (Cornwell et al 1999)), and its theoretical basis has come under occasional eriticism (Tan 1986... Lannes et al. 1997))."," The reasons for CLEAN's success were unclear for some time (Cornwell et al \cite{cornwell_1999}) ), and its theoretical basis has come under occasional criticism (Tan \cite{tan}, Lannes et al \cite{lannes}) )."
 Only relatively recently has it been shown to be related to compressive sampling (Candéss Wakin 2008)) and been given a sound theoretical underpinning., Only relatively recently has it been shown to be related to compressive sampling (Candèss Wakin \cite{compressive_sampling}) ) and been given a sound theoretical underpinning.
 In order for Hóggbom CLEAN to work. the convolution relation D=7*B must be valid.," In order for Höggbom CLEAN to work, the convolution relation $D=I \ast B$ must be valid."
 There are cases where this assumption breaks down however., There are cases where this assumption breaks down however.
 One of the most severe departures occurs in frequency synthesis with large fractional bandwidths., One of the most severe departures occurs in frequency synthesis with large fractional bandwidths.
 It is common for a single observational field to contain objects whose spectral indices differ by several tens of percent or more., It is common for a single observational field to contain objects whose spectral indices differ by several tens of percent or more.
 Differences in spectra of this order are unimportant if the observation fractional bandwidth is small but may significantly degrade the convolution assumption where the fractional bandwidth approaches 1., Differences in spectra of this order are unimportant if the observation fractional bandwidth is small but may significantly degrade the convolution assumption where the fractional bandwidth approaches 1.
 Conway et al (1990)) showed that. in the wide-band case emSeen which the dirty image D no longer well approximates a convolution of the sky brightness distribution 7. it was evertheless possible to express D as a sum over a relatively wonall number N of component images D;. each of which —nadividually obeys a convolution relation D;=J;«Bj.," Conway et al \cite{conway}) ) showed that, in the wide-band case in which the dirty image $D$ no longer well approximates a convolution of the sky brightness distribution $I$, it was nevertheless possible to express $D$ as a sum over a relatively small number $N$ of component images $D_i$, each of which individually obeys a convolution relation $D_i=I_i \ast B_i$."
 Conway et al arrived at this by expanding the nominal spectrum at each sky location ina Taylor series. each ΠΠ term of the series generating a respective “spectral dirty beam’ B;.," Conway et al arrived at this by expanding the nominal spectrum at each sky location in a Taylor series, each $i$ th term of the series generating a respective `spectral dirty beam' $B_i$."
 In this case each image /; is simply a sky map of the value of the ith Taylor coefficient., In this case each image $I_i$ is simply a sky map of the value of the $i$ th Taylor coefficient.
 Conway et al suggested a coordinate transform for better application of this technique to commonly-found power-law radio spectra. and presented an approximate method to solve for the coefficient images /; when N is restricted to 2.," Conway et al suggested a coordinate transform for better application of this technique to commonly-found power-law radio spectra, and presented an approximate method to solve for the coefficient images $I_i$ when $N$ is restricted to 2."
" This method consists of a 2-step sequential CLEAN and relies on the beams By and B, being approximately orthogonal.", This method consists of a 2-step sequential CLEAN and relies on the beams $B_0$ and $B_1$ being approximately orthogonal.
 Sault and Wieringa (1994)) retained the Taylor expansion but devised a new solution algorithm which can be thought of as a generalization of the Hóggbom CLEAN algorithm from its original “scalar context. in which a single image D ts iteratively deconvolved. to a new ‘vector’ context in which V images D; are deconvolved in parallel.," Sault and Wieringa \cite{sault_wieringa}) ) retained the Taylor expansion but devised a new solution algorithm which can be thought of as a generalization of the Höggbom CLEAN algorithm from its original `scalar' context, in which a single image $D$ is iteratively deconvolved, to a new `vector' context in which $N$ images $D_i$ are deconvolved in parallel."
 Orthogonality of the beams is no longer required (although the technique fails if 2 or more beams are identical)., Orthogonality of the beams is no longer required (although the technique fails if 2 or more beams are identical).
 Sault and Wieringa elaborated their theory as it applied to the N=2 case but. as the authors themselves suggest. the extensionto N»2 is not difficult.," Sault and Wieringa elaborated their theory as it applied to the $N=2$ case but, as the authors themselves suggest, the extensionto $N>2$ is not difficult."
 The CLEAN algorithm continues to be a subject of active development., The CLEAN algorithm continues to be a subject of active development.
 Recent work includes an extension of CLEAN to produce clean components with a rangeof sizes (Cornwell 2008 )). a modification to cleartomographic LISA images (Hayama et al 2006)). and an adaption of the algorithm to reconstruct RHESSI images of solar flares (Schwartz 2009)).," Recent work includes an extension of CLEAN to produce clean components with a rangeof sizes (Cornwell \cite{cornwell_2008}) ), a modification to cleantomographic LISA images (Hayama et al \cite{hayama}) ), and an adaption of the algorithm to reconstruct RHESSI images of solar flares (Schwartz \cite{schwartz}) )."
 The compressive-sampling formalism has recently also generated some promising new approaches (Wiaux et al 2000:: Li. Cornwell de Hoog 2011).," The compressive-sampling formalism has recently also generated some promising new approaches (Wiaux et al \cite{wiaux}; ; Li, Cornwell de Hoog \cite{li}) )."
For the idealized Newtonian hot-spot we assume that there is no lensing bv the black hole. the disk is completely optically thin. and the hot-spot is a point (the first two are physically. appropriate for spots at large orbital radii),"For the idealized Newtonian hot-spot we assume that there is no lensing by the black hole, the disk is completely optically thin, and the hot-spot is a point (the first two are physically appropriate for spots at large orbital radii)."
 In this case the hot spots orbit (which we assume to be cireular with radius r) is simply given by Kepler's law: where i is the orbital inclination. Algorys is (he mass of Ser A*. and © is an arbitraryphase.," In this case the hot spot's orbit (which we assume to be circular with radius $r$ ) is simply given by Kepler's law: where $i$ is the orbital inclination, $M_{\rm Sgr A*}$ is the mass of Sgr A*, and $\phi$ is an arbitraryphase."
 The image centroid. Xe is constructed by integrating (he source emission over a time T. which need not be small in comparison to the orbital period. and (hus the motion of the hot-spot will generally be important.," The image centroid, $\bmath{X}_C$ is constructed by integrating the source emission over a time $T$, which need not be small in comparison to the orbital period, and thus the motion of the hot-spot will generally be important."
 Explicitly. if we set the centroid of the disk emission to be al the origin. This will generally be a function of the initial phase of the orbit. ó. ancl the integration lime 7.," Explicitly, if we set the centroid of the disk emission to be at the origin, This will generally be a function of the initial phase of the orbit, $\phi$, and the integration time $T$."
 If we make the simplilving assumption that the hot-spot flix may be treated as roughly constant. then this reduces to and πώ:T) is the average spot position over some time T with initial orbital o.," If we make the simplifying assumption that the hot-spot flux may be treated as roughly constant, then this reduces to and $\overline{x}(\phi;T)$ is the average spot position over some time $T$ with initial orbital $\phi$ ."
 It is straightforward to show that, It is straightforward to show that
suggested by (LO84).,suggested by .
. In the theories of strong interaction quark bag models suppose that breaking of physical vacuum takes place inside hadrons., In the theories of strong interaction quark bag models suppose that breaking of physical vacuum takes place inside hadrons.
 If the hypothesis of the quark nmalter is (rue. then some compact objects idenüfied with neutron stus could actually be strange stars. built entirely of strange matter 1986).," If the hypothesis of the quark matter is true, then some compact objects identified with neutron stars could actually be strange stars, built entirely of strange matter ."
.. For a review of strange star properties. see (1998b).," For a review of strange star properties, see ."
. Most of the investigations of quark star properties have been done within the framework ol the so-called MIT bag model., Most of the investigations of quark star properties have been done within the framework of the so-called MIT bag model.
 Assuming that interactions of quarks and gluons are sullicientlv small. neglecting quark masses and supposing Chat quarks are confined to the bag volume (in the case of a bare strange star. the boundary of the bag coincides with stellar surface). the energy density pe? and pressure p of a quark-gluon plasma ave related. in the AUT bag model. by the equation of state (EOS) where 2 is the difference between the energy densitv of the perturbative and QCD vacuum (the bag constant).," Assuming that interactions of quarks and gluons are sufficiently small, neglecting quark masses and supposing that quarks are confined to the bag volume (in the case of a bare strange star, the boundary of the bag coincides with stellar surface), the energy density $\rho c^{2}$ and pressure $p$ of a quark-gluon plasma are related, in the MIT bag model, by the equation of state (EOS) where $B$ is the difference between the energy density of the perturbative and non-perturbative QCD vacuum (the bag constant)."
 Equation (1)) is essentially the equation of state of a gas of massless particles with corrections due to the QCD trace anomaly ancl perturbative interactions., Equation ) is essentially the equation of state of a gas of massless particles with corrections due to the QCD trace anomaly and perturbative interactions.
 More sophisticated investigations of quark-gluon interactions have shown that Eq. (1), More sophisticated investigations of quark-gluon interactions have shown that Eq. )
) represents a limiüng case of more general equations of state., represents a limiting case of more general equations of state.
 For example MIT bag models with massive strange quarks and lowest order QCD interactions lead (o some corrections terms in (he equation of state of quark matter., For example MIT bag models with massive strange quarks and lowest order QCD interactions lead to some corrections terms in the equation of state of quark matter.
 Models incorporating restoration of chiral quark masses al high densities and giving absolutelv stable strange matter can no longer be accurately described bv using Eq. (1))., Models incorporating restoration of chiral quark masses at high densities and giving absolutely stable strange matter can no longer be accurately described by using Eq. ).
 On the other hand in models in which quark interaction is described by an interquark potential originating Irom eluon exchange and by a densitwv dependent scalar potential which restores (he chiral svaumetry. at high densities 1993).. the equation of state P?=Pí(p) can be well approximated by a linear function in the energy density p2000)..," On the other hand in models in which quark interaction is described by an interquark potential originating from gluon exchange and by a density dependent scalar potential which restores the chiral symmetry at high densities , the equation of state $P=P\left( \rho \right) $ can be well approximated by a linear function in the energy density $\rho $."
 has studied the linear approximation of the equation of state. obtaining all parameters of the EOS as polvnonual functions of strange quark mass. QCD coupling constant aud bag constant.," has studied the linear approximation of the equation of state, obtaining all parameters of the EOS as polynomial functions of strange quark mass, QCD coupling constant and bag constant."
 A complete description of static strange stars has been obtained based on numerical integration of mass continuitv and TOV (lvdrostatic equilibrium) equations for different values of the bag constant1986)., A complete description of static strange stars has been obtained based on numerical integration of mass continuity and TOV (hydrostatic equilibrium) equations for different values of the bag constant.
". Using numerical methods the maximum gravitational mass AM,,,, and the maximum radius Ay,» of (he strange star. have been obtained. as a function of thebag constant. in the form"," Using numerical methods the maximum gravitational mass $M_{max}$ and the maximum radius $R_{max}$ of the strange star, have been obtained, as a function of thebag constant, in the form"
eoverned by the chirp mass M=(mq+ma)2Oneyh? (Peters and. Matthews. 1963).,"governed by the chirp mass $\chirp=(m_1+m_2)^{2/5}(m_1 m_2)^{3/5}$ (Peters and Matthews, 1963)."
 The waveform will depend on the individual masses of (he binary components mq ancl m» when the post Newtonian effects ave taken into account., The waveform will depend on the individual masses of the binary components $m_1$ and $m_2$ when the post Newtonian effects are taken into account.
 However. (he analvsis of the inspiral phase alone shall not sullice to determine if a binary contained a neutron star or a black hole without the prior knowledge of the neutron star maximum mass.," However, the analysis of the inspiral phase alone shall not suffice to determine if a binary contained a neutron star or a black hole without the prior knowledge of the neutron star maximum mass."
 A careful modeling of the signal may vield the individual masses of the objects. however. the chirp mass will be the primary observable for the compact object mergers (Cutler Flanagan 1994).," A careful modeling of the signal may yield the individual masses of the objects, however, the chirp mass will be the primary observable for the compact object mergers (Cutler Flanagan 1994)."
 We use theTrack population synthesis code (Belezvuski. Ikalogera 22002) to calculate the distributions of compact object binary masses. and we present these ealeulations in 2.," We use the population synthesis code (Belczynski, Kalogera 2002) to calculate the distributions of compact object binary masses, and we present these calculations in 2."
 In 3 we estimate the number of merger detections required to distinguish between different models of stellar binary evolution., In 3 we estimate the number of merger detections required to distinguish between different models of stellar binary evolution.
 Finally. 4 contains the conclusions and discussion.," Finally, 4 contains the conclusions and discussion."
 The StarTrack binary population synthesis code is described in detail in Belezvuski. Kalogera ((2002).," The StarTrack binary population synthesis code is described in detail in Belczynski, Kalogera (2002)."
 One of the important leatures of the code is the possibility to conduct parameter study of a given property of the population of binaries. Le. to estimate the dependence of the result on each of the parameters used to describe the stellar and binary evolution.," One of the important features of the code is the possibility to conduct parameter study of a given property of the population of binaries, i.e. to estimate the dependence of the result on each of the parameters used to describe the stellar and binary evolution."
 The models used in this paper are listed in Table 1., The models used in this paper are listed in Table 1.
 We first use the standard model A results to present the intrinsic distribution of the chirp masses., We first use the standard model A results to present the intrinsic distribution of the chirp masses.
 This is shown in Figure 1., This is shown in Figure \ref{chirpgal}.
", The distribution shows a clear peak at low chirp masses LAL.M<2M. which is due to the double neutron star svstems.", The distribution shows a clear peak at low chirp masses $1 M_\odot\chirp < 2 M_\odot$ which is due to the double neutron star systems.
 The mixed (DII-NS) svstems populate the intermediate region. while (he chirp masses of the DII-DII binaries extend up to above LOAL..," The mixed (BH-NS) systems populate the intermediate region, while the chirp masses of the BH-BH binaries extend up to above $10 M_\odot$."
 In order to estimate the observed distribution of the chirp masses of compact objects one has to take into account the sensitivity of the gravitational wave detectors (ο signals from mergers of different binaries., In order to estimate the observed distribution of the chirp masses of compact objects one has to take into account the sensitivity of the gravitational wave detectors to signals from mergers of different binaries.
 The calculations of the signal to noise ratio (Finn and Chernoll 1993. Bonazzola and Marck 1994. Flanagan and Hughes 1997) show that the sampling distance in the first approximation is a function of the chirp mass onlv: DxM5.," The calculations of the signal to noise ratio (Finn and Chernoff 1993, Bonazzola and Marck 1994, Flanagan and Hughes 1997) show that the sampling distance in the first approximation is a function of the chirp mass only: $D\propto \chirp^{5/6}$."
 The additional corrections due to limited sensitivity window of the detectors have been calculated bv Flanagan and Hughes (1997) and amount to approximately for the binaries with the total mass below L&Af. for the initial LIGO. and less for the advanced. LIGO.," The additional corrections due to limited sensitivity window of the detectors have been calculated by Flanagan and Hughes (1997) and amount to approximately for the binaries with the total mass below $18 M_\odot$ for the initial LIGO, and less for the advanced LIGO."
 In this paper we neelect these corrections., In this paper we neglect these corrections.
 The distribution of the expected observed chirp masses can be caleulated using Monte Carlo method., The distribution of the expected observed chirp masses can be calculated using Monte Carlo method.
 We assume that the Universe is uniformly filled with merging binaries. and lor each merger we estimate the signal to noise ratio in the detector.," We assume that the Universe is uniformly filled with merging binaries, and for each merger we estimate the signal to noise ratio in the detector."
 We model (he population of merging binaries assuming a continuous star formation rate., We model the population of merging binaries assuming a continuous star formation rate.
 The, The
According to cosmological ACDM simulations. large galaxies such às the Milkv Way were formed hierarchically.,"According to cosmological $\Lambda$ CDM simulations, large galaxies such as the Milky Way were formed hierarchically."
 Evidence of spatial and kinematical substructures in the Galactic halo resulting from this process has indeed been found as reviewed by Helmi (2008)) and Klement (2010))., Evidence of spatial and kinematical substructures in the Galactic halo resulting from this process has indeed been found as reviewed by Helmi \cite{helmi08}) ) and Klement \cite{klement10}) ).
" Elemental abundance ratios of halo stars may also be used to probe this formation process by so-called ""chemical tagging’ of the “building blocks’ (Freeman Bland-Hawthorn 2002)).", Elemental abundance ratios of halo stars may also be used to probe this formation process by so-called `chemical tagging' of the `building blocks' (Freeman Bland-Hawthorn \cite{freeman02}) ).
" The F. G. and K stars are of particular interest in this connection. because their atmospheric composition are likely to provide a ""fossil record of the composition of the gas from which the stars once were formed."," The F, G, and K stars are of particular interest in this connection, because their atmospheric composition are likely to provide a `fossil' record of the composition of the gas from which the stars once were formed."
 As foundby Korn et al. (2007)), As foundby Korn et al. \cite{korn07}) )
 from a comparison of dwarf and giant stars in the globular cluster 66397. the atmospheric abundances of heavy elements like Mg. Ca. Ti. and Fe in old dwarf stars may be somewhat decreased by diffusion processes. but the ratios between these elements are practically unchanged.," from a comparison of dwarf and giant stars in the globular cluster 6397, the atmospheric abundances of heavy elements like Mg, Ca, Ti, and Fe in old dwarf stars may be somewhat decreased by diffusion processes, but the ratios between these elements are practically unchanged."
 Most high-precision studies of abundance ratios 1n. stars are limited to the solar region of the Galaxy., Most high-precision studies of abundance ratios in stars are limited to the solar region of the Galaxy.
 This means that halo stars must be identified by their kinematics. e.g. by having space velocities with respect to the local standard of rest (LSR) well above the characteristic velocities of thin and thick-disk stars.," This means that halo stars must be identified by their kinematics, e.g. by having space velocities with respect to the local standard of rest (LSR) well above the characteristic velocities of thin and thick-disk stars."
 Several such studies have focused on possible correlationsbetween kinematics and|a/Fe]. where a refers to the abundance of a-capture elements like Mg. Si. Ca and Ti.," Several such studies have focused on possible correlationsbetween kinematics and, where $\alpha$ refers to the abundance of $\alpha$ -capture elements like Mg, Si, Ca and Ti."
 Fulbright (2002)) finds evidence that stars with high values of the total space velocity relative to the LSR. Vi>300. tend to have lower values of tthan stars with 150«Visa<300.," Fulbright \cite{fulbright02}) ) finds evidence that stars with high values of the total space velocity relative to the LSR, $\Vtotal > 300$, tend to have lower values of than stars with $150 < \Vtotal < 300$."
 Stephens Boesgaard (2002)) show that iis correlated with the apogalactic orbital distance (Ryyo) in the sense that the outermost stars have the lowest values of[a/Fe|., Stephens Boesgaard \cite{stephens02}) ) show that is correlated with the apogalactic orbital distance ) in the sense that the outermost stars have the lowest values of.
 Gratton et al. (2003)), Gratton et al. \cite{gratton03}) )
 divided their sample into two populations according to kinematics: {) a ‘dissipative’ component. which comprises thick-disk stars and prograde-rotating halo stars. and ii) an ‘accretion’ component consisting of retrograde-rotating halo stars.," divided their sample into two populations according to kinematics: $i)$ a `dissipative' component, which comprises thick-disk stars and prograde-rotating halo stars, and $ii)$ an `accretion' component consisting of retrograde-rotating halo stars."
 The accretion component has lower values and a larger scatter for tthan the dissipative component., The accretion component has lower values and a larger scatter for than the dissipative component.
 This has been confirmed by Jonsell et al. (2005))., This has been confirmed by Jonsell et al. \cite{jonsell05}) ).
 Furthermore. a recent study of Ishigaki et al. (2010))," Furthermore, a recent study of Ishigaki et al. \cite{ishigaki10}) )"
" shows that in the metallicity range —2<[Fe/H] -]. the rratio for stars reaching a maximum vertical distance Skkpe above or below the Galactic plane in their orbits tend to be about ddex lower than Που stars with Zi,< Skkpe."," shows that in the metallicity range $-2 < \feh < -1$ , the ratio for stars reaching a maximum vertical distance $\Zmax > 5$ kpc above or below the Galactic plane in their orbits tend to be about dex lower than for stars with $\Zmax < 5$ kpc."
 The a-elements are mainly produced during Type II supernovae ILD explosions of massive stars on a relatively short time scale. 107 years. whereas iron is also produced by Type la SNe Ila) on a longer time scale. i.e. from about 10° to more than 10° years (Maoz et al. 2010)).," The $\alpha$ -elements are mainly produced during Type II supernovae II) explosions of massive stars on a relatively short time scale, $\sim \! 10^7$ years, whereas iron is also produced by Type Ia SNe Ia) on a longer time scale, i.e. from about $10^8$ to more than $10^{9}$ years (Maoz et al. \cite{maoz10}) )."
 The differences in mmay therefore be explained in terms of differences in the star formation rate (SFR)., The differences in may therefore be explained in terms of differences in the star formation rate (SFR).
 The outer halo stars may originate from regions characterized by a relatively slow SFR with both Ha and SNelll contributing to the chemical evolution. whereas the inner stars come from regions with such a fast chemical evolution that only HI have contributed.," The outer halo stars may originate from regions characterized by a relatively slow SFR with both Ia and II contributing to the chemical evolution, whereas the inner stars come from regions with such a fast chemical evolution that only II have contributed."
 Differences in the initial mass function may. however. also play a role. because low-mass SNell produce lower values of tthan high-mass SNell (e.g. Kobayashi et al. 2006)).," Differences in the initial mass function may, however, also play a role, because low-mass SNeII produce lower values of than high-mass SNeII (e.g. Kobayashi et al. \cite{kobayashi06}) )."
 It is unclear from the works cited above if there is a dichotomy in the distribution of or a more continuous change in aas a function of oor Zax., It is unclear from the works cited above if there is a dichotomy in the distribution of or a more continuous change in as a function of or .
 In a recent study. the authors of the present paper (Nissen Schuster 2010.. hereafter NSIO) have. however. found evidence for the existence of two distinct," In a recent study, the authors of the present paper (Nissen Schuster \cite{nissen10}, , hereafter NS10) have, however, found evidence for the existence of two distinct"
Που a single 20-min scan and [for the full clataset.,for a single 20-min scan and for the full dataset.
 For a mean 19-12 flux. density of 177.0mniJ.. this corresponds to a 5m upper limit of 0.17 per cent on the lincar polarization.," For a mean 15-GHz flux density of mJy, this corresponds to a $5\sigma$ upper limit of 0.17 per cent on the linear polarization."
 At Ες. however. significant linear. polarization was detected. at. levels of up to 2.5mm.Js.," At GHz, however, significant linear polarization was detected, at levels of up to mJy."
 Phe level of linear. polarization initially decreases in line with the total intensity. but then rises during and after the Hare at ~ 21hh (Fig. 2)).," The level of linear polarization initially decreases in line with the total intensity, but then rises during and after the flare at $\sim21$ h (Fig. \ref{fig:polarization}) )."
 No circular »olarization. was detected at either frequency to a So Limit of, No circular polarization was detected at either frequency to a $5\sigma$ limit of.
 Linear polarization has occasionally. been detected. in oevious giant [lares of (νο N-3 (e.gSeaquistctal.1974: 1976). up to levels of ~20 per cent. with the ractional polarization increasing with frequency.," Linear polarization has occasionally been detected in previous giant flares of Cyg X-3 \citep[e.g][]{Sea74,Led76}, up to levels of $\sim20$ per cent, with the fractional polarization increasing with frequency."
 However. he majority of observations have shown no significant »obarization. with upper limits of a [few per cent (e.g.Draes&Alilev1972:οςBalick 1972). so there are relatively few datasets available for comparison.," However, the majority of observations have shown no significant polarization, with upper limits of a few per cent \citep[e.g.][]{Bra72,Hje72}, so there are relatively few datasets available for comparison."
 TPudoseetal.(2007). used. the European VLDE. Network (EWN) o resolve polarized emission. from. νο N-3. finding a maximum fractional polarization. of 25 per cent at the roundary where an ejected: knot was interacting with its environment.," \citet{Tud07} used the European VLBI Network (EVN) to resolve polarized emission from Cyg X-3, finding a maximum fractional polarization of 25 per cent at the boundary where an ejected knot was interacting with its environment."
 For the derived. rotation measure (RAL) of 1233 mm (Leddenctal.1976)... we would expec a rotation of 0.5 rrad at 15€GCGllz. and. O.0Grracd a 43€GGllz.," For the derived rotation measure (RM) of $-1233$ $^{-2}$ \citep{Led76}, we would expect a rotation of $-0.5$ rad at GHz, and $-0.06$ rad at GHz."
 This RAL was derived: during giant outbursts. when the emitting regions had. propagated further from the core than those we see during these small Daring events.," This RM was derived during giant outbursts, when the emitting regions had propagated further from the core than those we see during these small flaring events."
 | is likely therefore. that there is some extra Faraday rotation in the inner region. causing cepolarization at. the lower frequeney either due to averaging within the synthesisec beam of the observations or along the line of sight.," It is likely therefore, that there is some extra Faraday rotation in the inner region, causing depolarization at the lower frequency either due to averaging within the synthesised beam of the observations or along the line of sight."
 Llowever. since the theoretical maximum fractional. polarization. for an optically thick source is 12 per cent. as compared to το per cent for an optically thin source. we would in any case expect a lower degree of polarization at CiCillz. where the optical depth is greater.," However, since the theoretical maximum fractional polarization for an optically thick source is $\sim 12$ per cent, as compared to $\sim 70$ per cent for an optically thin source, we would in any case expect a lower degree of polarization at GHz, where the optical depth is greater."
 “Phe beam and line-ol-sight depolarization then reduce the observed. degree of polarization still further. below our detection threshold.," The beam and line-of-sight depolarization then reduce the observed degree of polarization still further, below our detection threshold."
 At GCGIIZ. the polarization position angle is stable to within 20 eexcept during the [lare ab — hh. which causes à significant. albeit temporary. rotation of the polarization position angle. while the cegree of linear. polarization increases slightlv.," At GHz, the polarization position angle is stable to within $\sim20$ except during the flare at $\sim21$ h, which causes a significant, albeit temporary, rotation of the polarization position angle, while the degree of linear polarization increases slightly."
 Lf this [lare is due to the formation of internal shocks in the Low (see Section 4)). the resulting compression of the magnetic field at the shock front. could be responsible for the rotation of the polarization position angle and the increase in polarized intensitv.," If this flare is due to the formation of internal shocks in the flow (see Section \ref{sec:modelling}) ), the resulting compression of the magnetic field at the shock front could be responsible for the rotation of the polarization position angle and the increase in polarized intensity."
 Bearing in mind the uncertainty in absolute position angle calibration. we see the mean EVPA. evolve from 585+107 at the start of the observations. through 76°+4 during the lave at 21hh. back to S4415 at the end of the observations (as shown in Fig. 2))," Bearing in mind the uncertainty in absolute position angle calibration, we see the mean EVPA evolve from $-58^{\circ}\pm10^{\circ}$ at the start of the observations, through $76^{\circ}\pm4^{\circ}$ during the flare at h, back to $-84^{\circ}\pm15^{\circ}$ at the end of the observations (as shown in Fig. \ref{fig:polarization}) )."
 For optically-thin emission. the measured EVRA should. be perpendicular o the magnetic field orientation.," For optically-thin emission, the measured EVPA should be perpendicular to the magnetic field orientation."
 Previous high-resolution observations (Mioduszewskietal.POOL:Aliller-Jonesοἱal.2004:Tudoseet2007). have shown a north-south jet axis with a position angle close to 0," Previous high-resolution observations \citep{Mio01,Mil04,Tud07} have shown a north-south jet axis with a position angle close to $0^{\circ}$."
 Thus a field xincipallv aligned along the jet axis should have an EEVIPA of ~907. while compression due to à shock viewed sidewavys should give an IEEVPA of ~07.," Thus a field principally aligned along the jet axis should have an EVPA of $\sim90^{\circ}$, while compression due to a shock viewed sideways should give an EVPA of $\sim0^{\circ}$."
 While the EVPA at the end. of the observations is consistent witha field. aligned ong the jet axis. there are several possible explanations or the intermediate. position angles scen earlier.," While the EVPA at the end of the observations is consistent witha field aligned along the jet axis, there are several possible explanations for the intermediate position angles seen earlier."
 In. the Fence of precession. either there is additional Faraday rotation beyond the 3° expected From the previously derived QM. the EVPA calibration is inaccurate. or we are seeing emission [from a superposition of emitting regions. some with a longitudinal field configuration (upstream of the shock ront). and some from the compressed. fields at the shock ront.," In the absence of precession, either there is additional Faraday rotation beyond the $3^{\circ}$ expected from the previously derived RM, the EVPA calibration is inaccurate, or we are seeing emission from a superposition of emitting regions, some with a longitudinal field configuration (upstream of the shock front), and some from the compressed fields at the shock front."
 Alternatively. if we are secing the shock front close to ace on. as has been proposed. for Cygnus X-3 (Lindforsetal. 2007).. owing to the small angle between the jet axis and the line of sight. (Mioduszewskietal.2001:Miller-Jonesetal. 2004)... we would not expect to see an increase in the degree of ordering of the field.," Alternatively, if we are seeing the shock front close to face on, as has been proposed for Cygnus X-3 \citep{Lin07}, owing to the small angle between the jet axis and the line of sight \citep{Mio01,Mil04}, we would not expect to see an increase in the degree of ordering of the field."
 In. the absence of data at different frequencies. or at higher spatial and time resolution. it is not possible to differentiate between. these competing explanations.," In the absence of data at different frequencies, or at higher spatial and time resolution, it is not possible to differentiate between these competing explanations."
 The rise in fractional polarization towards the end of the observation could. be explained. in part. by. the expansion of the emitting region following the laree fare at the start of the observations., The rise in fractional polarization towards the end of the observation could be explained in part by the expansion of the emitting region following the large flare at the start of the observations.
 As the source. becomes more optically thin. the theoretical maximum. fractional polarization increases.," As the source becomes more optically thin, the theoretical maximum fractional polarization increases."
 Furthermore. the density. of the stellar wind. of the companion star decreases as the gecta move outwards. reducing the effect. of Faraday depolarization in the wind.," Furthermore, the density of the stellar wind of the companion star decreases as the ejecta move outwards, reducing the effect of Faraday depolarization in the wind."
 Alternatively. a change in the relative dominance of a depolarized core ancl ejecta. as seen in NTIS 11748-288. (Brocksopp 2007).. could. be behind. the increase in. fractional polarization (consistent with the hypothesis of multiple emitting regions proposed to explain the observed. Αν).," Alternatively, a change in the relative dominance of a depolarized core and highly-polarized ejecta, as seen in XTE 1748-288 \citep{Bro07}, , could be behind the increase in fractional polarization (consistent with the hypothesis of multiple emitting regions proposed to explain the observed EVPAs)."
"ab fig,=2x 100m (2000 km) above the photosphere (Allen1973). to an arbitrary outer boundary at roy=300HR...","at $h_{\rm ch}=2\times 10^8$ cm (2000 km) above the photosphere \citep{aln73}
 to an arbitrary outer boundary at $r_{\rm out}=300 R_{\odot}$."
" We set four boundary conditions to construct a unique solution for a given input wave energv [αν £i. as follows: where F.,,,; in eq.(17)) is a maximum value of downward conductive flix in (he inner corona."," We set four boundary conditions to construct a unique solution for a given input wave energy flux $F_{{\rm w}, 0}$ as follows: where $F_{\rm c, max}$ in \ref{eq:bc3}) ) is a maximum value of downward conductive flux in the inner corona."
 The first condition denotes that wave energy flix must agree witli à given value when the waves start to dissipate., The first condition denotes that wave energy flux must agree with a given value when the waves start to dissipate.
 The second condition is also straightforward: the temperature has to coincide with the fixed value at the inner boundary., The second condition is also straightforward: the temperature has to coincide with the fixed value at the inner boundary.
 The (hird. condition is the requirement (hat the downward thermal conductive flux should become sulliciently small ab the upper chromosphere (T= 107). diminishing from its enormous value at the coronal base (P~ LON).," The third condition is the requirement that the downward thermal conductive flux should become sufficiently small at the upper chromosphere $T=10^4$ K), diminishing from its enormous value at the coronal base $T\sim 10^6$ K)."
 Practically. we continue calculations iteratively until Fora)/Fou<1% is satisfied.," Practically, we continue calculations iteratively until $F_{\rm c}(r_{\rm ch})/F_{\rm c, max}<1\%$ is satisfied."
 The fourth condition corresponds to an ordinary requirement that no heat is conducted inward [rom infinitv (5andboek&Leer1994)., The fourth condition corresponds to an ordinary requirement that no heat is conducted inward from infinity \citep{sl94}.
. Note that thanks to Che thirel condition. coronal base density. which is poorly determined [rom the observations. does nol have to be used as a boundary. condition.," Note that thanks to the third condition, coronal base density, which is poorly determined from the observations, does not have to be used as a boundary condition."
 As a result. the number of [ree parameters to be set in advance is reduced. (Ilamuner1932a.b;:Withbroe1938).," As a result, the number of free parameters to be set in advance is reduced \citep{hm82a,hm82b,wtb88}."
. The densitv at the coronal base or the (ransiüon region (TR) is calculated as an output: larger input. Fio increases downward F. in the lower corona. demanding larger density in the coronal base and TR to enhance radiative cooling to balance with the increased conductive heating.," The density at the coronal base or the transition region (TR) is calculated as an output; larger input $F_{{\rm w}, 0}$ increases downward $F_{\rm c}$ in the lower corona, demanding larger density in the coronal base and TR to enhance radiative cooling to balance with the increased conductive heating."
 For numerical integration of the momentum equation (eq.(9))) ancl (hie energy equation (eq.(11))). we respectively use e and an isothermal sound velocity. a. defined as To carry out the integration. the equations shown in the previous section need to be transformed into useful Forms.," For numerical integration of the momentum equation \ref{eq:eqm}) )) and the energy equation \ref{eq:egcns}) )), we respectively use $v$ and an isothermal sound velocity, $a$, defined as To carry out the integration, the equations shown in the previous section need to be transformed into useful forms."
 First. an expression for velocity gradient can be written from," First, an expression for velocity gradient can be written from"
"In the early decades of astrophysics, star clusters have been our main key to the understanding of stellar evolution.","In the early decades of astrophysics, star clusters have been our main key to the understanding of stellar evolution."
" While clusters continue to provide precious constraints on stellar physics, they are today studied in their own right and as tracers of the histories of galaxies."," While clusters continue to provide precious constraints on stellar physics, they are today studied in their own right and as tracers of the histories of galaxies."
" It has become clear that a significant fraction of star formation occurs in clusters, and that events such as interactinget galaxies can [Barton,trigger Geller,their formation (Harriset [Matteoal.|/2007)."," It has become clear that a significant fraction of star formation occurs in clusters, and that events such as interacting galaxies can trigger their formation \citep{Harris1991,Meurer1995, Barton2000, DiMatteo2007}."
". Questions have been raised& regarding the IMF in clusters in various environments, about the systematic trends in their colour distributions, about their lifetimes as gravitationally bound objects and about the initial and current cluster mass functions."," Questions have been raised regarding the IMF in clusters in various environments, about the systematic trends in their colour distributions, about their lifetimes as gravitationally bound objects and about the initial and current cluster mass functions."
 Resolved observations of individual stars remain the most precise way of investigating the nature of clusters and will be possible out to distances of 10 Mpc with future extremely large telescopes., Resolved observations of individual stars remain the most precise way of investigating the nature of clusters and will be possible out to distances of 10 Mpc with future extremely large telescopes.
" However measurements of the integrated light of unresolved star clusters reach far beyond this scale already today, and will remain the path of choice for the studies of large samples."," However measurements of the integrated light of unresolved star clusters reach far beyond this scale already today, and will remain the path of choice for the studies of large samples."
" All our studies of individual clusters and of cluster populations in galaxies rest on our ability to estimate their current ages, masses and metallicities, while accounting for extinction."," All our studies of individual clusters and of cluster populations in galaxies rest on our ability to estimate their current ages, masses and metallicities, while accounting for extinction."
 The standard method of analysis of integrated cluster light is based on the direct comparison of the observed colours with predictions frommodels., The standard method of analysis of integrated cluster light is based on the direct comparison of the observed colours with predictions from.
 These models predict fluxes with the assumption that each mass bin along the stellar mass function (SMF) is populated according to the average value given by this SMF., These models predict fluxes with the assumption that each mass bin along the stellar mass function (SMF) is populated according to the average value given by this SMF.
" Studies based on continuous population synthesis models have led to results that have a large impact on today's description of cluster ""demographics""."," Studies based on continuous population synthesis models have led to results that have a large impact on today's description of cluster “demographics""."
" For instance, it is now usually admitted that the current cluster mass function decreases with &mass as a 999;power[Bik etlaw with an index close to —2 and Fallthe debate on al]the cluster survival rate also rests on distributions obtained using continuous models[sky][2005)."," For instance, it is now usually admitted that the current cluster mass function decreases with mass as a power law with an index close to $-2$ \citep{
Zhang1999, Bik2003, Boutloukos2003} and the debate on the cluster survival rate also rests on distributions obtained using continuous models."
". The continuous approach has been coupled with statistical data analysis, for instance to provide the impression that including near-IR photometry (K band) solves the [PuziametallicityetaL 002;degeneracy[Anders foret aLclusters2004; Bridzius(GoudfrooijetaLD008)..al.200Tf"," The continuous approach has been coupled with statistical data analysis, for instance to provide the impression that including near-IR photometry (K band) solves the age-metallicity degeneracy for clusters \citep{Goudfrooij2001, Puzia2002, Anders2004,
Bridzius2008}."
 The continuous population synthesis models are strictly valid only in the limit of a stellar population containing an infinite number of stars., The continuous population synthesis models are strictly valid only in the limit of a stellar population containing an infinite number of stars.
" Real clusters, however, count a finite number of stars."," Real clusters, however, count a finite number of stars."
" Furthermore most of the light is provided by a very small number of bright stars, in particular in the near-IR."," Furthermore most of the light is provided by a very small number of bright stars, in particular in the near-IR."
 The so-calledfluctuations in the integrated photometric properties are the result of the random presence of these luminous stars., The so-called in the integrated photometric properties are the result of the random presence of these luminous stars.
" Some of these can be quantified using selected information provided by continuous population synthesis models (e.g.&(CervifioLuridrana[2004] 2006),, but others require the use of Popescu&Hanson|2009: PiskunoaI]2009)."," Some of these can be quantified using selected information provided by continuous population synthesis models \citep[e.g.][]{Lancon2000, Cervino2002, Cervino2004,
Cervino2006}, but others require the use of \citep{Barbaro1977, Girardi1993, Bruzual2002,Deveikis2008,Popescu2009,
Piskunov2009}."
". The predicted luminosity and colour distributions eidepend strongly on the total mass (or star number) in the cluster, and can be far from Gaussian even when the total mass exceeds 10° Mo."," The predicted luminosity and colour distributions depend strongly on the total mass (or star number) in the cluster, and can be far from Gaussian even when the total mass exceeds $10^5$ $_{\odot}$."
" The most probable colours are offset from those predicted by continuous population synthesis when masses are below 10* Mo, because the single most luminous star in such clusters will be more often on the main sequence than in the red giant phases of evolution."," The most probable colours are offset from those predicted by continuous population synthesis when masses are below $10^4$ $_{\odot}$, because the single most luminous star in such clusters will be more often on the main sequence than in the red giant phases of evolution."
" Attempts to describe the colour distributions analytically have made progress (e.g. [Luridiana|2006)., but are not yet easily applicable."," Attempts to describe the colour distributions analytically have made progress \citep[e.g.][]{Cervino2006}, but are not yet easily applicable."
 The present piece of work is based on discrete population synthesis., The present piece of work is based on discrete population synthesis.
" For the first time, we use the discrete models not only to predict colour distributions but to the energy distributions of clusters."," For the first time, we use the discrete models not only to predict colour distributions but to the energy distributions of clusters."
" We present a Bayesian approach to the probabilistic determination of age, mass and extinction, based on a large library of Monte-Carlo simulations of clusters."," We present a Bayesian approach to the probabilistic determination of age, mass and extinction, based on a large library of Monte-Carlo simulations of clusters."
that the lags of 100 5 were not present in the first oobservation of $5 OF16+714.,that the lags of $\gtrsim100$ s were not present in the first observation of S5 0716+714.
 However. we found that the 0.30.5 keV variations lag the 310 keV ones bv ~1000 s in the second oobservation.," However, we found that the 0.3–0.5 keV variations lag the 3–10 keV ones by $\sim 1000$ s in the second observation."
 We also [found a weak evidence that the lags increase with larger energy differences., We also found a weak evidence that the lags increase with larger energy differences.
 In al least one episode of the optical-UV light curves. the U band. variations might lag the N-rav. variations by ~2000 s. As far às we know. it is the first evidence [or a definite detection of soft lag in the X-ray. variability of 55 07164114 and possibly LBLs.," In at least one episode of the optical-UV light curves, the U band variations might lag the X-ray variations by $\sim 2000$ s. As far as we know, it is the first evidence for a definite detection of soft lag in the X-ray variability of S5 0716+714 and possibly LBLs."
 Interestingly. the soft lags ancl the related energy dependence of 55 07164714 are similar to what have been detected in HDBLs (e.g.. INataoka οἱ al.," Interestingly, the soft lags and the related energy dependence of S5 0716+714 are similar to what have been detected in HBLs (e.g., Kataoka et al."
 2000: Zhang et al., 2000; Zhang et al.
 2010)., 2010).
 The similarity suggests that the X-ray variations of the source have the same origin as those of 1119115. ie.. the variations of the svnehrotron tail.," The similarity suggests that the X-ray variations of the source have the same origin as those of HBLs, i.e., the variations of the synchrotron tail."
 Therefore. the hard X-ray. [hixes of the source are dominated bv the IC component. but the bard X-ray variations might be still controlled by the synchrotron tail. as already suggested by the energy dependent variability amplitude.," Therefore, the hard X-ray fluxes of the source are dominated by the IC component, but the hard X-ray variations might be still controlled by the synchrotron tail, as already suggested by the energy dependent variability amplitude."
" If. we assume that both the soft and harc X-ray variations are caused by the svuchrotron tail. the observed soft lags provide a way to constrain (he physical parameters ol the emitting region (e.g.. Zhang 2002). where 7,5 15 the observed soft lag (in second) between the low (£1) and high (£4) energy (in keV). 2 the source's redshilt. D the magnetic field (in G) and 9 the bulk Doppler factor of the emitting region."," If we assume that both the soft and hard X-ray variations are caused by the synchrotron tail, the observed soft lags provide a way to constrain the physical parameters of the emitting region (e.g., Zhang 2002), where $\tau_{\rm soft}$ is the observed soft lag (in second) between the low $E_{\rm l}$ ) and high $E_{\rm h}$ ) energy (in keV), $z$ the source's redshift, $B$ the magnetic field (in G) and $\delta$ the bulk Doppler factor of the emitting region."
 If adopting zi1000 s between the 0.30.5 and 310 keV. one gets Bol’ο2.56 G. During a model fit to the SED involving the first oobservation of S5 O716+714. Foschini et al. (," If adopting $\tau_{\rm soft} \sim 1000$ s between the 0.3–0.5 and 3–10 keV, one gets $B\delta^{1/3} \sim 2.56$ G. During a model fit to the SED involving the first observation of S5 0716+714, Foschini et al. ("
2006) assumed 5=3 G and 9=16.7. Le. BOY’=τον G. A larger Do! implies a smaller Το. qualitatively consistent with the soft lags of no larger than 100 s claimed by FEOQG.,"2006) assumed $B=3$ G and $\delta =16.7$, i.e., $B\delta^{1/3}
= 7.67$ G. A larger $B\delta^{1/3}$ implies a smaller $\tau_{\rm
soft}$, qualitatively consistent with the soft lags of no larger than 100 s claimed by FE06."
 The unprecedented. high signal-to-noise ratio PN observation allows us to disentangle the svncehrotron ancl IC components in the X-ray. bad of $5 07164714 and to svnchronously study (he variations of the (wo components on timescales of hours., The unprecedented high signal-to-noise ratio PN observation allows us to disentangle the synchrotron and IC components in the X-ray band of S5 0716+714 and to synchronously study the variations of the two components on timescales of hours.
 The results. obtained with the divisions of the observation bv individual episodes aud by count rate levels. are consistent with each other.," The results, obtained with the divisions of the observation by individual episodes and by count rate levels, are consistent with each other."
 The svnchrotron. photon indices are constrained in a limited range of D~2.5—2.7. which are (vpical of IIBLs in the hard N-rays (e.g.. PINS 2155304: Zhang 2008).," The synchrotron photon indices are constrained in a limited range of $\Gamma \sim 2.5-2.7$, which are typical of HBLs in the hard X-rays (e.g., PKS 2155–304: Zhang 2008)."
 The IC photon indices show relatively large changes (D~0.9— 1.4)., The IC photon indices show relatively large changes $\Gamma \sim 0.9-1.4$ ).
 Due to large uncertainties. (he svnchrotron photon indices might be consistent with each other within the error bars. which could be also true for the IC photon indices.," Due to large uncertainties, the synchrotron photon indices might be consistent with each other within the error bars, which could be also true for the IC photon indices."
 Although the svnchrotron and IC photon indices do not correlate with the total fluxes. the svncehrotron spectra appear to harden with higher svnchrotron fluxes. and the IC spectra seem to soften," Although the synchrotron and IC photon indices do not correlate with the total fluxes, the synchrotron spectra appear to harden with higher synchrotron fluxes, and the IC spectra seem to soften"
" 20804 thy ο”) y (2s Mase, 2054 thy Pon=0.( tha dey thydsey)=0.( where ez>=οσαiyo is. the square ofJ the equilibrium“ye"" sound speed and an over-dot denotes a time clerivative.","_x - v_y 	+ ik_x = 0, _y + (2- v_x 	+ ik_y = 0, + c_s^2 (i k_x v_x 	+ i k_y v_y) = 0, where $c_s^2 = \gamma P_0/\Sigma_0$ is the square of the equilibrium sound speed and an over-dot denotes a time derivative."
 The above system of equations adiits four linearly-uilependent. solutious., The above system of equations admits four linearly-independent solutions.
 Two of these are the nou-vortical sliwaves (solutions for which the perturbed. poteutial vorticity is zero). which in the absence of sell-gravity can be solved [or exactly.," Two of these are the non-vortical shwaves (solutions for which the perturbed potential vorticity is zero), which in the absence of self-gravity can be solved for exactly."
 The remaining two solutions are the vortical suwaves., The remaining two solutions are the vortical shwaves.
" When &,—0 the latter reduce to the zero-frequency modes of the axisymmetric vers]on o. equations (??)) through (??)).", When $k_y \rightarrow 0$ the latter reduce to the zero-frequency modes of the axisymmetric version of equations \ref{LIN1s}) ) through \ref{LIN4s}) ).
 One of these (the eutropy mode) remains uuchauged iu uonaxisyumetry (in a frame comoving with the shear)., One of these (the entropy mode) remains unchanged in nonaxisymmetry (in a frame comoving with the shear).
 There is thus only one noutrivial vortical sluvave iL the unstratifiecd shearing sheet., There is thus only one nontrivial vortical shwave in the unstratified shearing sheet.
 In tle limit of tightly-wound sliwaves l|m hy). the uonvortical aud. vortical shwaves are compressive aud ---iucoumpressive. respectively.," In the limit of tightly-wound shwaves $|k_x| \gg k_y$ ), the nonvortical and vortical shwaves are compressive and incompressive, respectively."
" In the slort-wavelength limit. (Ημ>1. where ds te vertical scale height). the compressive aud i1compressive solutions remain well separated at all times. but for Hy,SO(1) there is mixing between them near fy=0 as au iucompressive slwave stears from leadiug to With the uilerstaudiug that the distinction between compressive sliwaves aud incompressive sliwaves as separate solutious is not valid for all time when Hbi,SOC). we generally choose to employ these te‘us over the more general but less intuitive lernis ""uco-vortical aud “vortical.”"," In the short-wavelength limit $H k_y \gg 1$, where $H \equiv c_s/\Omega$ is the vertical scale height), the compressive and incompressive solutions remain well separated at all times, but for $H k_y \lesssim O(1)$ there is mixing between them near $k_x = 0$ as an incompressive shwave shears from leading to With the understanding that the distinction between compressive shwaves and incompressive shwaves as separate solutions is not valid for all time when $H k_y \lesssim O(1)$, we generally choose to employ these terms over the more general but less intuitive terms “non-vortical” and “vortical.”"
" Based upon the above considerations. it is convenient to study the vortical shwave in the short-waveleugth. low-frequency (0,< y) limit."," Based upon the above considerations, it is convenient to study the vortical shwave in the short-wavelength, low-frequency $\partial_t \ll 
c_s k_y$ ) limit."
 This is equivalent to working iu the Boussinesq which in the uustratifiec shearing sheet amounts to asstuning incompressible flow., This is equivalent to working in the Boussinesq which in the unstratified shearing sheet amounts to assuming incompressible flow.
 In this liuit. equation (?2)) is replaced with kde d keydty = 0.," In this limit, equation \ref{LIN1s}) ) is replaced with k_x v_x + k_y v_y = 0."
 In this liuit. equation (?2)) is replaced with kde d keydty = 0.(," In this limit, equation \ref{LIN1s}) ) is replaced with k_x v_x + k_y v_y = 0."
 In this liuit. equation (?2)) is replaced with kde d keydty = 0.(1," In this limit, equation \ref{LIN1s}) ) is replaced with k_x v_x + k_y v_y = 0."
 In this liuit. equation (?2)) is replaced with kde d keydty = 0.(12," In this limit, equation \ref{LIN1s}) ) is replaced with k_x v_x + k_y v_y = 0."
 In this liuit. equation (?2)) is replaced with kde d keydty = 0.(12)," In this limit, equation \ref{LIN1s}) ) is replaced with k_x v_x + k_y v_y = 0."
We have studied the astrometric aspects of microlensing by simultaneously including the FS and FL effects.,We have studied the astrometric aspects of microlensing by simultaneously including the FS and FL effects.
" Our results show that the astrometric signal is underestimated or overestimated by assuming PL or PS, respectively."," Our results show that the astrometric signal is underestimated or overestimated by assuming PL or PS, respectively."
" While the FS effect is prominent when the lens transits the surface of the source, the FL effect is revealed when the lens is very close to the source, which would be in the self-lensing regime."," While the FS effect is prominent when the lens transits the surface of the source, the FL effect is revealed when the lens is very close to the source, which would be in the self-lensing regime."
" In the context of the self-lensing scenario, where a background star is lensed by a foreground star, the light contribution from the lens is in general not negligible."," In the context of the self-lensing scenario, where a background star is lensed by a foreground star, the light contribution from the lens is in general not negligible."
" We thus consider the luminous-lens scenario, which attenuates the signal of the centroidal displacement."," We thus consider the luminous-lens scenario, which attenuates the signal of the centroidal displacement."
" Astrometric trajectories with a source located in the Galactic bulge, SMC, and M31 are discussed, which show that Og of halo-lensing events is at least one order of magnitude larger than that of self-lensing in SMC and M31."," Astrometric trajectories with a source located in the Galactic bulge, SMC, and M31 are discussed, which show that $\AERR$ of halo-lensing events is at least one order of magnitude larger than that of self-lensing in SMC and M31."
" Our results also indicate that the finiteness of the lens is more likely to be revealed in the self-lensing scenario towards distant source located in Magellanic Clouds or M31, although it is very difficult to distinguish between PL and FL with current instruments."," Our results also indicate that the finiteness of the lens is more likely to be revealed in the self-lensing scenario towards distant source located in Magellanic Clouds or M31, although it is very difficult to distinguish between PL and FL with current instruments."
 We are very grateful to the anonymous referee for the useful comments., We are very grateful to the anonymous referee for the useful comments.
 This work was supported by the DFG cluster of excellence ‘Origin and Structure of the Universe’, This work was supported by the DFG cluster of excellence `Origin and Structure of the Universe' (www.universe-cluster.de).
Semi-analytic models and numerical simulations predict that the first astrophysical objects formed at redshifts ο 20.,Semi-analytic models and numerical simulations predict that the first astrophysical objects formed at redshifts $z\gsim$ 20.
" They form out of metal-free gas with inefficient fragmentation. thus numerical simulations predict that these first objects could be very massive (AL,z100A. ). so-called “Population IIT. (Pop TIT) stars (222)."," They form out of metal-free gas with inefficient fragmentation, thus numerical simulations predict that these first objects could be very massive $M_\ast\gsim 100 \Msun$ ), so-called 'Population III' (Pop III) stars \citep{ABN02, BCL02, YOH08}."
 These stars would be short-lived. with lifetimes of ~3 Myr. but could produce ten times more ionizing photons per baryon than regular Population IT stars (22)...," These stars would be short-lived, with lifetimes of $\sim 3$ Myr, but could produce ten times more ionizing photons per baryon than regular Population II stars \citep{Schaerer02, Schaerer03}."
 Also. such star formation could have been somewhat synchronized. given that these early objects formed in highly biased and clustered environments.," Also, such star formation could have been somewhat synchronized, given that these early objects formed in highly biased and clustered environments."
 Hence. Pop III stars could have a signiticant impact on subsequent generations of objects.," Hence, Pop III stars could have a significant impact on subsequent generations of objects."
" Indeed Pop III stars can strongly affect the progress. of cosmological reionization (2222)... although their contribution is virtually unknown from first principles and is generally applied to simulations with a ""tuning knob"" approach."," Indeed Pop III stars can strongly affect the progress of cosmological reionization \citep{HH03, Cen03a, Cen03b, WL03_postWMAP}, although their contribution is virtually unknown from first principles and is generally applied to simulations with a “tuning knob"" approach."
 Thus ionizing photons from the first stars and lack thereof have been used to explain both early (e.g. 2») and fairly late (e.g. 25) reionization. depending on which scenario was favored when the works were published.," Thus ionizing photons from the first stars and lack thereof have been used to explain both early (e.g. \citealt{Cen03b}) ) and fairly late (e.g. \citealt{HB06}) ) reionization, depending on which scenario was favored when the works were published."
connected. with the quasars.,connected with the quasars.
 I. is possible that the if the absorbing material is intrinsic to these quasars. then it could be similar in origin to the high ionisation absorbers observed in more nearby AGN.," It is possible that the if the absorbing material is intrinsic to these quasars, then it could be similar in origin to the high ionisation absorbers observed in more nearby AGN."
 The most substantial source of soft X-ray absorption [reom intervening matter would. probably be from damped nance absorption systems. however the number densitv of hese dense systems is reported. το be quite low (COFIaherv Jakobsen 1997).," The most substantial source of soft X-ray absorption from intervening matter would probably be from damped $\alpha$ absorption systems, however the number density of these dense systems is reported to be quite low (O'Flaherty Jakobsen 1997)."
 A more detailed account of the possie origins of this X-ray absorption will not be discussed. furher in this paper. but can be found in the literature (e.g. Evis 1998. Cappl 1997. Reeves LOOT. Elvis 1994).," A more detailed account of the possible origins of this X-ray absorption will not be discussed further in this paper, but can be found in the literature (e.g. Elvis 1998, Cappi 1997, Reeves 1997, Elvis 1994)."
 Llowever observations using the superior low energy throughput provided by NMM ane Chandra are needed to confirm this trend. and το provide further clues as to the possible causes.," However observations using the superior low energy throughput provided by XMM and Chandra are needed to confirm this trend, and to provide further clues as to the possible causes."
 We now summarize the findings of this paper., We now summarize the findings of this paper.
 In particular. comparison is crawn to our earlier paper (Reeves 1997 or ROT). which contained a smaller sunple of quasars (24 objects compared with the current sample size of 62).," In particular, comparison is drawn to our earlier paper (Reeves 1997 or R97), which contained a smaller sample of quasars (24 objects compared with the current sample size of 62)."
 Firstly we confirm the following main results [rom the LOT paper: Llowever we have seen several new collects ane correlations in this paper. that were not reported in our previous ROT sample: So how can we place all these observations [acts into a general scheme for quasars.," Firstly we confirm the following main results from the R97 paper:- However we have seen several new effects and correlations in this paper, that were not reported in our previous R97 sample:- So how can we place all these observations facts into a general scheme for quasars."
 Firstly. the differences between racdio-Ioud and raclio-quiet quasars seem relatively straightforward., Firstly the differences between radio-loud and radio-quiet quasars seem relatively straightforward.
 In the radio-loud quasars. a strong Doppler boosted emission component. [from the relativistic jet can account lor the higher luminosities. the generally. Latter Xoray spectra as well as the diminished. iron Ix line and rellection component in these objects.," In the radio-loud quasars, a strong Doppler boosted emission component from the relativistic jet can account for the higher luminosities, the generally flatter X-ray spectra as well as the diminished iron K line and reflection component in these objects."
 One question of real, One question of real
1900)... galaxies.,", galaxies."
 Although a number of great results [from studies with the Sloan Digital Sky Survey (SDSS:Yorketal.2000) in the last vears have appeared. these works have not studied the astrophysical parameters targeted here.," Although a number of great results from studies with the Sloan Digital Sky Survey \citep[SDSS; ][]{Yorketal00} in the last years have appeared, these works have not studied the astrophysical parameters targeted here."
 As a part of a European Virtual Astronomical Infrastructure for Data Access (Euro-VO AIDA) research. initiative. we have undertaken a comprehensive analysis of the scale length in clisk galaxies using an unprecedentedly large sample of disk galaxies.," As a part of a European Virtual Astronomical Infrastructure for Data Access (Euro-VO AIDA) research initiative, we have undertaken a comprehensive analysis of the scale length in disk galaxies using an unprecedentedly large sample of disk galaxies."
 We have used the Virtual Observatory (VO) tools to retrieve data in all (wv. gor. 7. and 2) bands from the sixth SDSS major data release (DRG:AXdelman-MceCarthyetal.2008) which includes imaging catalogues. spectra. ancl redshifts freely available.," We have used the Virtual Observatory (VO) tools to retrieve data in all $u$, $g$, $r$, $i$, and $z$ ) bands from the sixth SDSS major data release \citep[DR6; ][]{AMcetal08} which includes imaging catalogues, spectra, and redshifts freely available."
 We use the catalogue (Patureletal.2003) to retrieve morphological classification information about our sample galaxies. and those with types defined as Sa or later are hereafter. refereed to as disk .galaxies (distribution of both samples are presented in Fig. 1)).," We use the catalogue \citep{Petal03} to retrieve morphological classification information about our sample galaxies, and those with types defined as Sa or later are hereafter refereed to as disk galaxies (distribution of both samples are presented in Fig. \ref{fig:radec}) )."
 In the present paper. we present the data retrieval and analysis method. used to automatically derive the scale lengths for a sample of disk galaxies which contains 56096 objects (described in section 2)). and after rigourous tests described. in section 3.. we [ind that a subset of. 30374 of these can be called. reliable. following these criteria.," In the present paper, we present the data retrieval and analysis method used to automatically derive the scale lengths for a sample of disk galaxies which contains 56096 objects (described in section \ref{sec:sdsssample}) ), and after rigourous tests described in section \ref{sec:scalelengths}, we find that a subset of 30374 of these can be called reliable following these criteria."
 The scale. lengths. presented here relate only to the. disk components. anc we have tried to avoid the regions that could be dominated by the bulge component. in order to avoid complications related to the uncertainties of bulge-disk decomposition procedure (asdemonstratedin.e.g...Inapen&vander," The scale lengths presented here relate only to the disk components, and we have tried to avoid the regions that could be dominated by the bulge component, in order to avoid complications related to the uncertainties of bulge-disk decomposition procedure \citep[as demonstrated in, e.g., ][]{KvK91}."
Wruit 1991).. In. section. d. we present. the first results based. on our unprecedentedly large sample of galaxies and finally discuss their implications in section 5.., In section \ref{sec:results} we present the first results based on our unprecedentedly large sample of galaxies and finally discuss their implications in section \ref{sec:discussion}.
 The DRG provides imaging catalogues. spectra. ancl redshifts for the third and final data release of SDSS-LL. an extension of the original SDSS consisting of three sub. projects: The Legacy Survey. the Sloan Extension for Galactic Understanding ancl Exploration. ancl a Supernova survey.," The DR6 provides imaging catalogues, spectra, and redshifts for the third and final data release of SDSS-II, an extension of the original SDSS consisting of three sub projects: The Legacy Survey, the Sloan Extension for Galactic Understanding and Exploration, and a Supernova survey."
 The SDSS Catalogue Archive Server Jobs allow for a sample selection based. on a number of useful morphological and spectroscopic parameters provided for all objects., The SDSS Catalogue Archive Server Jobs allow for a sample selection based on a number of useful morphological and spectroscopic parameters provided for all objects.
 We use these parameters and make a first selection of the entire SDSS DAG sample., We use these parameters and make a first selection of the entire SDSS DR6 sample.
 Various VO methods were investigated to perform the download of the SDSS images. and the SkyView was chosen for this task.," Various VO methods were investigated to perform the download of the SDSS images, and the SkyView was chosen for this task."
 This service has the advantage of being able to create fits cut outs centred at à given sky coordinate and with a given size., This service has the advantage of being able to create fits cut outs centred at a given sky coordinate and with a given size.
 Moreover. SkyView is able to re-scale the image backgrounds to the same level. hence correcting for background level dillerences between the SDSS tiles.," Moreover, SkyView is able to re-scale the image backgrounds to the same level, hence correcting for background level differences between the SDSS tiles."
 The image size is an important. parameter to achieve a reliable sky subtraction which is necessary to clerive realistic scale lengths. thus we require that the images cover an area at least three times the size of cach galaxy.," The image size is an important parameter to achieve a reliable sky subtraction which is necessary to derive realistic scale lengths, thus we require that the images cover an area at least three times the size of each galaxy."
 To optimise the cata handling and keep low data transfer time [rom SkyView. we chose a constant image size of 900.000 pixels to be sampled. for all galaxies. still being able to achieve a reliable sky subtraction.," To optimise the data handling and keep low data transfer time from SkyView, we chose a constant image size of $900 \times 900$ pixels to be sampled for all galaxies, still being able to achieve a reliable sky subtraction."
 With these specifications. the image size is 3.2 MD with the typical download time of 16 seconds per image.," With these specifications, the image size is 3.2 MB with the typical download time of 16 seconds per image."
 This also includes the time that SkyView spends cutting. mosaicing. and re-scaling images.," This also includes the time that SkyView spends cutting, mosaicing, and re-scaling images."
 Our first selection. criteria uses SDSS parameters to ensure that: This first set of criteria leaves us with a total of 05735 ealaxiecs., Our first selection criteria uses SDSS parameters to ensure that: This first set of criteria leaves us with a total of 95735 galaxies.
 We use the LEDA services to retrieve a numeric Llubble classification parameter 7 [or the galaxies in our sample (more on this in section 4.1))., We use the LEDA services to retrieve a numeric Hubble classification parameter $T$ for the galaxies in our sample (more on this in section \ref{sec:spirals}) ).
 We first. download he entire LEDA catalogue. which we cross-correlate with he SDSS sample using and only select. the galaxies which. in LEDA. are classified. as spiral galaxics (Le. 1.P< 5).," We first download the entire LEDA catalogue, which we cross-correlate with the SDSS sample using and only select the galaxies which, in LEDA, are classified as spiral galaxies (i.e., $1 \le T \le 8$ )."
 A total of 56096 Sa-Scl (oe. Z between 1 ancl 8) spiral galaxies (see Fig.," A total of 56096 Sa-Sd (i.e., $T$ between 1 and 8) spiral galaxies (see Fig."
" 1) were found. for which SDSS ""Lg.Ενἐν z-band images were downloaded."," 1) were found, for which SDSS $u, g, r, i,
z$ -band images were downloaded."
 In section 4.1. we urther discuss whether all these galaxies are well-classified disk or spiral galaxies.," In section \ref{sec:spirals}, we further discuss whether all these galaxies are well-classified disk or spiral galaxies."
 In Fig. 2..," In Fig. \ref{fig:design},"
 we show the distribution of some kev xwanmeters retrieved from the SDSS and. LEDA database., we show the distribution of some key parameters retrieved from the SDSS and LEDA database.
 This figure shows that the cdillerent sample selection stages do not introduce any biases in our sample., This figure shows that the different sample selection stages do not introduce any biases in our sample.
 Ht should be noted hat. at this stage. we are unable to determine whether the ealaxies in our sample are isolated or cisturbed systems. as his information is not provided by any of the catalogues we jte used.," It should be noted that, at this stage, we are unable to determine whether the galaxies in our sample are isolated or disturbed systems, as this information is not provided by any of the catalogues we have used."
 We make this distinction using the asymmetry xuwameter described in Schadeetal. (1995)..., We make this distinction using the asymmetry parameter described in \citet{Schadeetal95}. .
unreliability of the stellarity parameter at fainter optical magnitudes 1s not biasing this result.,unreliability of the stellarity parameter at fainter optical magnitudes is not biasing this result.
 We note that for the sources with low X-ray counts. the X-ray flux may have been over-estimated due to Eddington bias.," We note that for the sources with low X-ray counts, the X-ray flux may have been over-estimated due to Eddington bias."
 This effect 1s discussed in Kenter et al. (, This effect is discussed in Kenter et al. (
"2005) and may mean that f;/f£,<10 fora fraction of our optically obscured AGN sample.",2005) and may mean that $f_x/f_o<10$ for a fraction of our optically obscured AGN sample.
 The X-ray hardness ratio is a useful quantity to approximate the spectral shape of the X-ray emission in cases for which the counting statistics are insufficient to determine the shape of the X-ray spectrum and the redshift distribution is unknown., The X-ray hardness ratio is a useful quantity to approximate the spectral shape of the X-ray emission in cases for which the counting statistics are insufficient to determine the shape of the X-ray spectrum and the redshift distribution is unknown.
" The X-ray hardness ratio is defined as: where C, and C, are the counts in the hard (2-7 keV) and soft (0.5-2 keV) bands respectively (Kenter at al.", The X-ray hardness ratio is defined as: where $C_h$ and $C_s$ are the counts in the hard (2-7 keV) and soft (0.5-2 keV) bands respectively (Kenter at al.
 2005)., 2005).
 In general. the intrinsic X-ray spectral slope is not thought to vary significantly from source to source.," In general, the intrinsic X-ray spectral slope is not thought to vary significantly from source to source."
 We therefore assume that changes in the hardness ratio can be attributed to a variation in the absorption column of gas which preferentially attenuates the soft X-rays., We therefore assume that changes in the hardness ratio can be attributed to a variation in the absorption column of gas which preferentially attenuates the soft X-rays.
 The hardness ratio spans the range —-| (only soft band counts) to 1 (only hard band counts: appropriate for obscured AGN in which the soft photons are absorbed by an intervening column of gas)., The hardness ratio spans the range $-$ 1 (only soft band counts) to 1 (only hard band counts: appropriate for obscured AGN in which the soft photons are absorbed by an intervening column of gas).
" In Figure 13.. we show the distribution of X-ray hardness ratio for the optically obscured AGN sample (f/f, 10) compared to that of the (renormalized) main AGN sample (O.1<fFf, <10)."," In Figure \ref{fig:hrhist}, we show the distribution of X-ray hardness ratio for the optically obscured AGN sample $f_x/f_o>$ 10) compared to that of the (renormalized) main AGN sample $<f_x/f_o<$ 10)."
" The optically obscured AGN sample exhibits a slightly broader distribution of X-ray hardness ratios whereas the —f/f, <10 sample is more dominated by soft X-ray sources.", The optically obscured AGN sample exhibits a slightly broader distribution of X-ray hardness ratios whereas the $<f_x/f_o<$ 10 sample is more dominated by soft X-ray sources.
 A KS test rules out the hypothesis that the two hardness ratio distributions are drawn from the same underlying distributions at a probability of 799., A KS test rules out the hypothesis that the two hardness ratio distributions are drawn from the same underlying distributions at a probability of $>$.
956... This suggests that AGN that are obscured by dust in the optical are also more likely to be obscured in the soft X-ray by gas., This suggests that AGN that are obscured by dust in the optical are also more likely to be obscured in the soft X-ray by gas.
 However. there is clearly no direct correspondence betwee! these quantities: many of the optically obscured X-ray sources exhibit hardness ratios consistent with them being unobscurec in the X-ray and conversely. many optically unobscured X-ray sources exhibit hardness ratios consistent with them being obscured 1n the X-ray.," However, there is clearly no direct correspondence between these quantities: many of the optically obscured X-ray sources exhibit hardness ratios consistent with them being unobscured in the X-ray and conversely, many optically unobscured X-ray sources exhibit hardness ratios consistent with them being obscured in the X-ray."
 This is a similar result to that found for a smaller number of sources by ?.., This is a similar result to that found for a smaller number of sources by \citet{wan04}.
" The sources with Ff. 10 are generally optically fainter (by selection) anc may be at a higher redshift than those with O.1<f/f, «10.", The sources with $f_x/f_o>$ 10 are generally optically fainter (by selection) and may be at a higher redshift than those with $<f_x/f_o<$ 10.
 If this is the case. the negative k-correction in the X-ray will result in similarly obscured sources exhibiting softer," If this is the case, the negative k-correction in the X-ray will result in similarly obscured sources exhibiting softer"
Class 0 objects are protostellar objects at the beginning of the accretion phase. where the bulk of the final stellar material has not been assembled yet (André et al. 1993)).,"Class 0 objects are protostellar objects at the beginning of the accretion phase, where the bulk of the final stellar material has not been assembled yet (André et al. \cite{andre:etal93}) )."
 These objects have been detected towards various regions with known ongoing star formation (André et al.1993... 1999;;," These objects have been detected towards various regions with known ongoing star formation (André et al.\cite{andre:etal93}, \cite{andre:etal99};"
 Kauffmann et al. 2005)., Kauffmann et al. \cite{kauffmann:etal05}) ).
 Molecular cirrus clouds at high Galactic latitudes so far show only very little star-formation activity. no embedded infrared sources had been found in the IRAS faint source catalog (Magnani et al. 1995)).," Molecular cirrus clouds at high Galactic latitudes so far show only very little star-formation activity, no embedded infrared sources had been found in the IRAS faint source catalog (Magnani et al. \cite{magnani:etal95}) )."
 Nevertheless. the existence of a gravitationally bound core in a cirrus clouds (Heithausen et al. 2002.. 2008)).," Nevertheless, the existence of a gravitationally bound core in a cirrus clouds (Heithausen et al. \cite{heithausen:etal02}, \cite{heithausen:etal08}) ),"
which shows evidence of inward motion (Heithausen 1999)). shows that these clouds harbor potential sites of star formation.,"which shows evidence of inward motion (Heithausen \cite{heithausen99}) ), shows that these clouds harbor potential sites of star formation."
 One candidate for possible star formation in such clouds is the dark cloud 11457. located at southern Galactic latitudes below the Taurus region.," One candidate for possible star formation in such clouds is the dark cloud 1457, located at southern Galactic latitudes below the Taurus region."
 This cloud is puzzling for two reasons: To shed light on its star-forming capability. we obtained a bolometer map of the densest part of the 11457 / 112 cloud at a wavelength of 1.2mmm with the IRAM 30mm telescope.," This cloud is puzzling for two reasons: To shed light on its star-forming capability, we obtained a bolometer map of the densest part of the 1457 / 12 cloud at a wavelength of mm with the IRAM m telescope."
" The selected region was found to be the most intense portion in the high angular resolution data of the ""CO and ""CO (1—ο. Q>1) and (3.—2) transitions obtained by Zimmermann (1993))."," The selected region was found to be the most intense portion in the high angular resolution data of the $^{12}$ CO and $^{13}$ CO $(1\to 0)$, $(2\to 1)$, and $(3\to 2)$ transitions obtained by Zimmermann \cite{zimmermann93}) )."
 Supplementary information on the dynamical state of the core comes from CS (2—1) observations obtained with the IRAM 30m telescope. too.," Supplementary information on the dynamical state of the core comes from CS $(2\to1)$ observations obtained with the IRAM 30m telescope, too."
 We discuss a possible association of two continuum point sources detected with the VLA at a wavelength of 3.6em and with Spitzer Space Telescope at 244m. 11457 Was observed in. 2000 and 2001 im the dust continuum at. l.2mm using the Max-Planck-Millimeter-Bolometerκά MAMBO at the IRAM 30m radio telescope on Pico Veleta. Spain.," We discuss a possible association of two continuum point sources detected with the VLA at a wavelength of 3.6cm and with Spitzer Space Telescope at $\mu$ m. 1457 was observed in 2000 and 2001 in the dust continuum at 1.2mm using the Max-Planck-Millimeter-Bolometer MAMBO at the IRAM 30m radio telescope on Pico Veleta, Spain."
 MAMBO is sensitive to emission between 210 ane GGHz. with an effective frequency of GGHz for steep thermal spectra.," MAMBO is sensitive to emission between 210 and GHz, with an effective frequency of GHz for steep thermal spectra."
" The observations were taken in double-beam on-the-fly mode. 1.e.. chopping the secondary mirror in azimuth by 50” to 70” at 2Hz and seanning the sky m azimuth at a speed of 4” to 5""57!. then moving in elevation by 4""."," The observations were taken in double-beam on-the-fly mode, i.e., chopping the secondary mirror in azimuth by $''$ to $''$ at 2Hz and scanning the sky in azimuth at a speed of $''$ to $'' s^{-1}$, then moving in elevation by $4''$."
 The maps were taken under variable winter conditions with line-of-sight opacities between 0.2 and 0.7., The maps were taken under variable winter conditions with line-of-sight opacities between 0.2 and 0.7.
 The effective beam FWHM is 11., The effective beam FWHM is $''$.
 The most intense dust positions ir 11457 were oberved in the CS (2>D) transition in October 2002 and in the N:H(1—0) transition in. August 2004 with the IRAM 30m radiotelescope., The most intense dust positions in 1457 were oberved in the CS $(2\to1)$ transition in October 2002 and in the $_2$ $^+ (1\to 0)$ transition in August 2004 with the IRAM 30m radiotelescope.
 Data were taken in single-position on-off-mode., Data were taken in single-position on-off-mode.
 We obtained small maps with 20” spacing betwee individual positions., We obtained small maps with $''$ spacing between individual positions.
 We used the VESPA autocorrelatior spectrometer with a velocity resolution set to 0.062 km s! at 93GHz and 0.060 km s! at 98GHz., We used the VESPA autocorrelation spectrometer with a velocity resolution set to 0.062 km $^{-1}$ at 93GHz and 0.060 km $^{-1}$ at 98GHz.
" The angular resolutior of the 30m telescope at 93GHz and 98 GHz is 27"" and 26"". respectively."," The angular resolution of the 30m telescope at 93GHz and 98 GHz is $''$ and $''$, respectively."
 The main beam efficiency μμ=0.8., The main beam efficiency $\eta_{mb}=0.8$.
 Additionally. we used public data from the VLA and Spitzer archives.," Additionally, we used public data from the VLA and Spitzer archives."
 In the VLA archive we found two thus far unpublished maps centered on about the same position as our bolometer map., In the VLA archive we found two thus far unpublished maps centered on about the same position as our bolometer map.
" The data were observed on March 6 and 25. 199] at a wavelength of 3.55cem with an angular resolution of 10""."," The data were observed on March 6 and 25, 1991 at a wavelength of cm with an angular resolution of $10''$."
 The primary beam of the VLA has a full width at half maximum of 574 and covers all our bolometer sources (s. Fig. 3))., The primary beam of the VLA has a full width at half maximum of $5.\!'4$ and covers all our bolometer sources (s. Fig. \ref{vla}) ).
 The noise level in the VLA map is 0.03 mJy/beam., The noise level in the VLA map is 0.03 mJy/beam.
 Part of the region was also observed with MIPS (Rieke et al. 2004)), Part of the region was also observed with MIPS (Rieke et al. \cite{rieke:etal04}) )
 onboard the Spitzer Space Telescope (Werner et al. 2004)), onboard the Spitzer Space Telescope (Werner et al. \cite{werner:etal04}) )
 at a wavelength of jm. The data were observed on September 23. 2007.," at a wavelength of $\mu$ m. The data were observed on September 23, 2007."
" The angular resolution of the MIPS instrument at that wavelength is about 6"". map spacing is 2.5""."," The angular resolution of the MIPS instrument at that wavelength is about $''$ , map spacing is $''$ ."
 The data taken from the Spitzer archive were already calibrated using, The data taken from the Spitzer archive were already calibrated using
unstable Πο burning precursor. aud the cooling of ashes of IC burning just below gp. dominate the cmission.,"unstable $^4$ He burning precursor, and the cooling of ashes of $^{12}$ C burning just below $y_{\rm He}$, dominate the emission."
 If unstable Πο burning is triggered. it results in a Type I burst that beeius to rise zLs after the peak. of the shock breakout flash.," If unstable $^4$ He burning is triggered, it results in a Type I burst that begins to rise $\approx 1\trm{ s}$ after the peak of the shock breakout flash."
 During the first ~Is of tHe burning. a convective zone develops aud moves outward from yy. to lower pressure (line 3 in Figure 9)).," During the first $\approx1\trm{ s}$ of $^4$ He burning, a convective zone develops and moves outward from $y_{\rm He}$ to lower pressure (line 3 in Figure \ref{fig:TyMol2}) )."
" The enerev generation rate €5, is so high during the initial stages of burning that the timescale for the convective zone to erow 1s muuch shorter than the thermal timescale at the couvective-radiative interface.", The energy generation rate $\epsilon_{3\alpha}$ is so high during the initial stages of burning that the timescale for the convective zone to grow is much shorter than the thermal timescale at the convective-radiative interface.
 As a result. the outer radiative lavers do not ect heated duiug the first second of Πο burning and the cooling wave continues to propagate inwards through the outer shock-heated lavers.," As a result, the outer radiative layers do not get heated during the first second of $^4$ He burning and the cooling wave continues to propagate inwards through the outer shock-heated layers."
 A steep temperature gradient develops at the couvective-radiative interface due to the large compositional contrast between the Ποσο matter that is burning aud the outer H-aich material (Weinbergetal. 2006b)., A steep temperature gradient develops at the convective-radiative interface due to the large compositional contrast between the $^4$ He-rich matter that is burning and the outer H-rich material \citep{Weinberg:06}.
. Eventually. the thermal time at the interface becomes shorter than the convective erowth time aud the convective zoue slowly retreats back to the base gp (lines {1.5.6.7} in Figure 9)).," Eventually, the thermal time at the interface becomes shorter than the convective growth time and the convective zone slowly retreats back to the base $y_{\rm He}$ (lines $\{4,5,6,7\}$ in Figure \ref{fig:TyMol2}) )."
 The radiative layers then finally start to heat wp aud the dux beeins to rise. reaching a peak luminosity after several secoucls.," The radiative layers then finally start to heat up and the flux begins to rise, reaching a peak luminosity after several seconds."
 The peak huuiuositv of the Type I burst increases as ο> Land the amount of Te available to burn iucreases (Figure 10))., The peak luminosity of the Type I burst increases as $\phi \rightarrow 1$ and the amount of $^4$ He available to burn increases (Figure \ref{fig:lightcurve2}) ).
 After the peak. the ΠΠο lavers cool aud the thermal wave coutinues its inward march. eradually penetrating the ashes of PC buruius (lines 18.9.LO} iu Figure 9)).," After the peak, the H/He layers cool and the thermal wave continues its inward march, gradually penetrating the ashes of $^{12}$ C burning (lines $\{8,9,10\}$ in Figure \ref{fig:TyMol2}) )."
 huportantle oven if unstable Πο burniug is uot triggered.ao the cooling ashes of °C burning just below Wye results in a flux rise ancl decay that is simular in appearance to that of a Type E burst. though with a peak flux that is smaller by a factor of &2.," Importantly, even if unstable $^4$ He burning is not triggered, the cooling ashes of $^{12}$ C burning just below $y_{\rm He}$ results in a flux rise and decay that is similar in appearance to that of a Type I burst, though with a peak flux that is smaller by a factor of $\approx 2$."
 This can be seen iu he o=0.5 curve of Figure 10.. a case in which the shock ‘ails to trigger unstable Πο burning (also compare lines D.1.5} of Figure δ with lines {3.1.5.6.7} of Figure 9)).," This can be seen in the $\phi=0.5$ curve of Figure \ref{fig:lightcurve2}, a case in which the shock fails to trigger unstable $^4$ He burning (also compare lines $\{3,4,5\}$ of Figure \ref{fig:TyMol1} with lines $\{3, 4,5,6,7\}$ of Figure \ref{fig:TyMol2}) )."
" Iu this case. the Type Llike rise aud decay is powered by he hot ashes of °C burning. although nuclear nnme of H and ‘Ile also contribute. increasing the o)eak flux by zLO%920% (see the bottom panel of Figure 7:: this may be an underestimate since we ouly account for stable IT burning via the rp-process up to 2181), burst."," In this case, the Type I-like rise and decay is powered by the hot ashes of $^{12}$ C burning, although nuclear burning of H and $^4$ He also contribute, increasing the peak flux by $\approx10\%-20\%$ (see the bottom panel of Figure \ref{fig:stableburning}; this may be an underestimate since we only account for stable H burning via the rp-process up to $^{24}$ Si)."
 Tn three of the five cases in which a precursor was seen. the peak flux of the precursor was smaller than the brightest of the svstenis ordinary bursts by a factor of 1.53 (the exceptions are the precursor from IU 30. which looked like au ordinary burst. aud 1U 69 which was brighter than au ordinary burst: Iuulkers2001:Crmumineetal. 20063).," In three of the five cases in which a precursor was seen, the peak flux of the precursor was smaller than the brightest of the system's ordinary bursts by a factor of $1.5-2$ (the exceptions are the precursor from 4U 1820-30, which looked like an ordinary burst, and 4U 1254-69 which was brighter than an ordinary burst; \citealt{Kuulkers:04, Cumming:06}) )."
 It is thus possible that some of the observed precursors are not the result of unstable Ie burnine., It is thus possible that some of the observed precursors are not the result of unstable $^4$ He burning.
 For fz20s. the light curve is powered bw the continued inward propagation of the cooling wave mto the ashes beneath yy. (see. Figure &.. lines {6.7.8.9}).," For $t\ga 20\trm{ s}$, the light curve is powered by the continued inward propagation of the cooling wave into the ashes beneath $y_{\rm He}$ (see Figure \ref{fig:TyMol1}, lines $\{6,7,8,9\}$ )."
 As Cunauiug&Macbeth(2001). showed. the light curve decays as a broken power-law. with the break marking the tine when the cooling wave first reaches g4: the deeper yp is. the later the break (Figure 11)).," As \citet{Cumming:04b} showed, the light curve decays as a broken power-law, with the break marking the time when the cooling wave first reaches $y_b$; the deeper $y_b$ is, the later the break (Figure \ref{fig:lightcurve1}) )."
" Our final temperature profile Ty(4) iu the detonation region Yer<yo<yp, ds onlv sliehtlv steeper than the profile iu the isobaric deflagration region due<8cLdae.", Our final temperature profile $T_f(y)$ in the detonation region $y_{\rm det} < y < y_b$ is only slightly steeper than the profile in the isobaric deflagration region $y_{\rm He} < y < y_{\rm det}$.
 Thus. our computed light curves at times £z100s are similar to those of Cunuuiug&Macbeth(2001).. who assumed au isobaric burn throughout.," Thus, our computed light curves at times $t \ga 100\trm{ s}$ are similar to those of \citet{Cumming:04b}, who assumed an isobaric burn throughout."
 Ilere. as in previous studies. the calculated slope of the pre-break light curve (fz100 1000s) i steeper than the observed slope (see c.g. Figure 6 of Cineetal. 20063).," Here, as in previous studies, the calculated slope of the pre-break light curve $t\approx 100-1000 \trm{ s}$ ) is steeper than the observed slope (see e.g., Figure 6 of \citealt{Cumming:06}) )."
 The emissiou during this epoch is powered by the cooling of the deflagration-heated reelon pucogcZYaer and in order to better match he observations. Tj(g) di this region must be steeper han we have assunued.," The emission during this epoch is powered by the cooling of the deflagration-heated region $y_{\rm He} < y < y_{\rm det}$, and in order to better match the observations, $T_f(y)$ in this region must be steeper than we have assumed."
 DIudeed. since the deflagration is ikelv to be highly turbulent due to the strong eravity and the hour-long convective churning that precedes he runaway. our assuniption of a complete isobaric mn nav not be correct.," Indeed, since the deflagration is likely to be highly turbulent due to the strong gravity and the hour-long convective churning that precedes the runaway, our assumption of a complete isobaric burn may not be correct."
 The deflagration iav mstead cave behind increasinely large pockets of unburned fuel as it accelerates down the density eradicut. thereby vielding the steepened temperature eradieut implied by he observed lisht curve.," The deflagration may instead leave behind increasingly large pockets of unburned fuel as it accelerates down the density gradient, thereby yielding the steepened temperature gradient implied by the observed light curve."
 We modeled the hydrodynamic IC combustion wave that forms diving a superburst as a detonation., We modeled the hydrodynamic $^{12}$ C combustion wave that forms during a superburst as a detonation.
 We found that the detonation propagates through the deepest lavers and drives a shock wave that steepeus as it travels upward., We found that the detonation propagates through the deepest layers and drives a shock wave that steepens as it travels upward.
 Upon reaching the I/We laver. the shock is sufficiently strong that it triggersao unstable Πο burning if," Upon reaching the H/He layer, the shock is sufficiently strong that it triggers unstable $^4$ He burning if"
We are now in the position to quantitatively auswer the original question: How fast would Bolt really have run. if he hadu't celebrated the last 2 seconds?,"We are now in the position to quantitatively answer the original question: How fast would Bolt really have run, if he hadn't celebrated the last 2 seconds?"
 To make this projection. we consider the following two scenarios: The justification of scenario Lis obvious. as Bolt outran Thompson between baud 8 seconds.," To make this projection, we consider the following two scenarios: The justification of scenario 1 is obvious, as Bolt outran Thompson between 4 and 8 seconds."
 The justification of scenario 2 dis more speculative. as it is difficult to quantity exactly stronger Bolt was.," The justification of scenario 2 is more speculative, as it is difficult to quantify exactly stronger Bolt was."
 Still looking at the acceleration profiles in Figure 2.. aud noting that Bolt traditionally was considered a 200 meter specialist. a value of 0.5 u/s? seenis fairly realistic.," Still, looking at the acceleration profiles in Figure \ref{fig:evolution}, and noting that Bolt traditionally was considered a 200 meter specialist, a value of 0.5 $^2$ seems fairly realistic."
 Then. for cach scenario we compute a new trajectory for Bolt bv choosing initial conditions. sys(Ssoc} and ο=c(8sec) and an acceleration profile as described above.," Then, for each scenario we compute a new trajectory for Bolt by choosing initial conditions, $s_0 = s(8 \,\textrm{sec})$ and $v_0 =
v(8\,\textrm{sec})$, and an acceleration profile as described above."
 The computation of these trajectories are performed by simply integrating Equation 1l with respect to time. Iu Figure 3. we compare the projected trajectories. s) (dashed red line shows scenario 1. dotted red line shows scenario 2). with the actual trajectory. s(t) (solid red line).," The computation of these trajectories are performed by simply integrating Equation \ref{eq:v_and_a} with respect to time, In Figure \ref{fig:end_run} we compare the projected trajectories, $\hat{s}(t)$ (dashed red line shows scenario 1, dotted red line shows scenario 2), with the actual trajectory, $s(t)$ (solid red line)."
 For comparison. Thompson's trajectory is indicated by a solid blue line.," For comparison, Thompson's trajectory is indicated by a solid blue line."
 The projected new world record is the time for which s(t} equals LOO meter., The projected new world record is the time for which $\hat{s}(t)$ equals 100 meter.
 Includiug statistical errors estimated by Moute Carlo simulations as described iu Section 2.. we fiud that the new world record would be 9.61£0.0L seconds in scenario d. and 9.55+0.01 in scenario 2.," Including statistical errors estimated by Monte Carlo simulations as described in Section \ref{sec:method}, , we find that the new world record would be $9.61\pm0.04$ seconds in scenario 1, and $9.55\pm0.04$ in scenario 2."
 Clen Mills. Usain Bolt’s coach. suggestedCoco that the world record could have been 9.52 secouds if Bolt had not danced along the track in Beijing for the last 20 meters.," Glen Mills, Usain Bolt's coach, suggested that the world record could have been 9.52 seconds if Bolt had not danced along the track in Beijing for the last 20 meters."
 According to our calculations. that secms like an good. but perhaps shehltly optimustic. estimate: Depending ou assuniptious about Bolt’s acceleration at the end of the race. we fud that his tine would have been somewhere between 9.55 aud 9.61 seconds. with a statistical error of EO.1 seconds.," According to our calculations, that seems like an good, but perhaps slightly optimistic, estimate: Depending on assumptions about Bolt's acceleration at the end of the race, we find that his time would have been somewhere between 9.55 and 9.61 seconds, with a statistical error of $\pm0.04$ seconds."
 Clearly. the unucertaiuties due to the assuniptious about the acceleration are comparable to or larger than the statistical uncertainties.," Clearly, the uncertainties due to the assumptions about the acceleration are comparable to or larger than the statistical uncertainties."
 Therefore. 9.52 seconds does by no means seem to be out of reach.," Therefore, 9.52 seconds does by no means seem to be out of reach."
 Iun Figure | we show an illustration of how such a record would compare to the actual world record of 9.69 seconds. relative to the rest of the field: The left version of Bolt shows lis actual position at ~9.5 secouds. while the right version indicates his position in the new scenarios.," In Figure \ref{fig:manipulated} we show an illustration of how such a record would compare to the actual world record of 9.69 seconds, relative to the rest of the field: The left version of Bolt shows his actual position at $\sim9.5$ seconds, while the right version indicates his position in the new scenarios."
 Of course. there are potential several systematics errors involved in these calculations.," Of course, there are potential several systematics errors involved in these calculations."
 For iustauce. it is inupossible to know for sure whether Usain might have been fired at the cud. which of course would increase the world record bevoud our estimates.," For instance, it is impossible to know for sure whether Usain might have been tired at the end, which of course would increase the world record beyond our estimates."
 On the other haud. judgiug frou his facial expressions as he crossed the finishine line. this doesut inuuediatelv strike us as a very plausible hypothesis.," On the other hand, judging from his facial expressions as he crossed the finishing line, this doesn't immediately strike us as a very plausible hypothesis."
 Another issue to consider is the wind., Another issue to consider is the wind.
 It is eenerallv agreed that a tail wind speed of Lo iu/s improves a 100 ineter time by 0.05 seconds (Aliweika2000)., It is generally agreed that a tail wind speed of 1 m/s improves a 100 meter time by 0.05 seconds \citep{mureika:2000}.
 Further. for IAAF (International Association of AthleticsFederations) to acknowledge a given run as a record," Further, for IAAF (International Association of AthleticsFederations) to acknowledge a given run as a record"
"the brightness temperature leads to In the frequency regime where the emission is dominated by free-free, Equations and show that the flux Εν of a (2)homogeneous source typically grows as v? in the optically thick regime (small and falls off like v~°-! in the optically thin regime frequencies)(large frequencies).","the brightness temperature leads to In the frequency regime where the emission is dominated by free-free, Equations and show that the flux $F_\nu$ of a homogeneous source typically grows as $\nu^2$ in the optically thick regime (small frequencies) and falls off like $\nu^{-0.1}$ in the optically thin regime (large frequencies)."
" For typical emission measures of EM=10? pc cm-$, the transition frequency for which T— lis 14£5.3 GHz."," For typical emission measures of $\mathrm{EM} = 10^8\,$ pc $^{-6}$, the transition frequency for which $\tau = 1$ is $\nu_\mathrm{t} \approx 5.3\,$ GHz."
" This turnover frequency depends only weakly on the emission measure, 14οςEM,"," This turnover frequency depends only weakly on the emission measure, $\nu_\mathrm{t} \propto \mathrm{EM}^{0.48}$."
" However, many observed SEDs of rregions do not behave in this simple manner but show anomalous scaling exponents (Lizano2008)."," However, many observed SEDs of regions do not behave in this simple manner but show anomalous scaling exponents \citep{lizano08}."
". In particular, there are SEDs that grow linearly with v over a frequency interval that can be as large as the whole VLA range (see Table 1))."," In particular, there are SEDs that grow linearly with $\nu$ over a frequency interval that can be as large as the whole VLA range (see Table \ref{vla}) )."
" Such anomalous scaling exponents can be reproduced by rregion models with ionized gradientsH (PanagiaKetoetal. 2008),, hierarchical clumps (Ignacewell2004) and additional dust emission (Rudolphetal. 2008)."," Such anomalous scaling exponents can be reproduced by region models with ionized gradients \citep{panfel75,olnon75,francoetal00,avaletal06,ketoetal08}, hierarchical clumps \citep{ignachur04} and additional dust emission \citep{rudetal90,pratetal92,beuthetal04,ketoetal08}."
. The data from our numerical simulations is certainly more realistic than the simple ionized gradient or hierarchical clump models., The data from our numerical simulations is certainly more realistic than the simple ionized gradient or hierarchical clump models.
 Therefore it is reassuring that we can reproduce the typical shapes of observed SEDs., Therefore it is reassuring that we can reproduce the typical shapes of observed SEDs.
 We show two examples in Figure 13.., We show two examples in Figure \ref{seds}.
 The left-hand SED shows the full coverage from VLA over ALMA to IRAS frequencies for the rregion of Figures 10 and 12.., The left-hand SED shows the full coverage from VLA over ALMA to IRAS frequencies for the region of Figures \ref{multifreq} and \ref{almapic}.
 The SEDs show the typical shape of those reported in Wood&Churchwell(1989)., The SEDs show the typical shape of those reported in \citet{woodchurch89}.
". As already noted above, the free-free radiation is by far dominant up to ALMA frequencies."," As already noted above, the free-free radiation is by far dominant up to ALMA frequencies."
" For IRAS frequencies, however, dust emission is the dominant process."," For IRAS frequencies, however, dust emission is the dominant process."
 The right-hand SED shows another example that more nicely illustrates the typical scaling behavior expected for the free-free emission., The right-hand SED shows another example that more nicely illustrates the typical scaling behavior expected for the free-free emission.
" The dotted lines grow οςν2, the dashed lines fall off οςv9."," The dotted lines grow $\propto \nu^2$, the dashed lines fall off $\propto \nu^{-0.1}$."
 One example of an SED with an anomalous scaling xri (solid line) is shown in Figure 14.., One example of an SED with an anomalous scaling $\propto \nu^1$ (solid line) is shown in Figure \ref{sedsabn}.
" The result can be a combined effect of density gradients in the ionized gas as well as shadowing by clumps, just as in the simpler analytical models."," The result can be a combined effect of density gradients in the ionized gas as well as shadowing by clumps, just as in the simpler analytical models."
" We do not need any dust to produce anomalous exponents, and in fact in our particular model, the dust emission is so low that it is unable to produce such gradients."," We do not need any dust to produce anomalous exponents, and in fact in our particular model, the dust emission is so low that it is unable to produce such gradients."
" Hence, our simulations reproduce both the general shape of rregions SEDs as well as the observed anomalous scaling behavior entirely with the computed distribution of ionized gas."," Hence, our simulations reproduce both the general shape of regions SEDs as well as the observed anomalous scaling behavior entirely with the computed distribution of ionized gas."
 We have presented synthetic continuum observations of the rregions formed in our collapse simulations of massive star formation described in Paper I. We find that the rregions are highly variable in size and shape as long as the massive star continues accreting., We have presented synthetic continuum observations of the regions formed in our collapse simulations of massive star formation described in Paper I. We find that the regions are highly variable in size and shape as long as the massive star continues accreting.
 Dense filaments in the gravitationally unstable accretion flow irregularly shield ionizing radiation from the central star., Dense filaments in the gravitationally unstable accretion flow irregularly shield ionizing radiation from the central star.
 This shielding effect leads to a large-scale (~5000 AU) flickering of the," This shielding effect leads to a large-scale $\sim 5000\,$ AU) flickering of the"
"where £,=dE1.",where $\dot{E}_c \equiv dE_c/dt$.
 Lhe rate of the cut-olf energy increase is à well-defined. quantity and the time scale (2.3) ias à clear physical interpretation., The rate of the cut-off energy increase is a well-defined quantity and the time scale (2.3) has a clear physical interpretation.
 The above definition does not require any limit for the energy. gains of the individual xwticles ancl all possible correlations are automatically included. here., The above definition does not require any limit for the energy gains of the individual particles and all possible correlations are automatically included here.
 From the meaning of the clefinition (2.3) it ollows that οι is somewhat shorter than the respective scale at the same energy for later times. required. for the respective part of the spectrum to become a pure power-law (cf.," From the meaning of the definition (2.3) it follows that $T_{acc}^{(c)}$ is somewhat shorter than the respective scale at the same energy for later times, required for the respective part of the spectrum to become a pure power-law (cf."
 Ostrowski Schlickeiser. 1996)., Ostrowski Schlickeiser 1996).
 One should. also note hat in relativistic shocks the time scale depends on the reference frame we use for its measurement., One should also note that in relativistic shocks the time scale depends on the reference frame we use for its measurement.
 In the present xiper the acceleration time scales are given in the respective normal shock rest frame., In the present paper the acceleration time scales are given in the respective normal shock rest frame.
 However. the applied time units reife (seo below) are defined with the use of the upstream gvration time.," However, the applied time units $r_{e,1}/c$ (see below) are defined with the use of the upstream gyration time."
 We derive estimates of τι basing on the Monte Carlo particle simulations involving the small amplitude momentum scattering. model. of Ostrowski (1991). and 2500 particles in each run of the code., We derive estimates of $T_{acc}^{(c)}$ basing on the Monte Carlo particle simulations involving the small amplitude momentum scattering model of Ostrowski (1991) and 2500 particles in each run of the code.
 In the present considerations we discuss the role of the mean magnetic field. configuration and the amount of particle scattering., In the present considerations we discuss the role of the mean magnetic field configuration and the amount of particle scattering.
 In order to avoid the additional cllects of the varving shock compression due to the presence of different magnetic field configurations we take the field as à trace one. without any dynamical effects. on the plasma Dow.," In order to avoid the additional effects of the varying shock compression due to the presence of different magnetic field configurations we take the field as a trace one, without any dynamical effects on the plasma flow."
 The shock compression. as seen in the normal shock rest frame. +. is derived. from. the approximate formulae. presented: by lleavens Drury (1989).," The shock compression, as seen in the normal shock rest frame, $r$, is derived from the approximate formulae presented by Heavens Drury (1989)."
 For illustration of the results. in the present. paper we consider the shock waves propagating in the cold electron-proton plasma.," For illustration of the results, in the present paper we consider the shock waves propagating in the cold electron-proton plasma."
" For the magnetic field D, taken in the upstream plasma rest frame and inclined at the angle ey with respect to the shock normal we derive its downstream value and inclination. D» and es. with the use of jump conditions presented. for relativistic shocks by c.g. ApplCamenzind (LOSS): £2 =re.feoand the Lorentz factors 5; Livi/U? G = 1. 2)."," For the magnetic field $B_1$ taken in the upstream plasma rest frame and inclined at the angle $\psi_1$ with respect to the shock normal we derive its downstream value and inclination, $B_2$ and $\psi_2$, with the use of jump conditions presented for relativistic shocks by e.g. ApplCamenzind (1988): where $R = r \, \gamma_1 / \gamma_2$ and the Lorentz factors $\gamma_i
\equiv 1/\sqrt{1-U_i^2}$ $i$ = $1$, $2$ )."
 These formulae ave valid for both sub- and super-Iuminal magnetic field configurations., These formulae are valid for both sub- and super-luminal magnetic field configurations.
 We model particle trajectory perturbations by introducing small-angle random momentum scattering along the mean-Llield trajectory., We model particle trajectory perturbations by introducing small-angle random momentum scattering along the mean-field trajectory.
 The. particle momentum scattering distribution is uniform within a cone wide at AQ (<< 1). along the original momentum cirection.," The particle momentum scattering distribution is uniform within a cone wide at $\Delta \Omega$ $<< 1$ ), along the original momentum direction."
 The present simulations use a constant value of AQ=0.173 (= 10°).," The present simulations use a constant value of $\Delta \Omega = 0.173$ $ =
10^\circ$ )."
 Scattering events are at discrete instants. equally. spaced in time as measured in the units of the respective Πο 0 — 1.2).," Scattering events are at discrete instants, equally spaced in time as measured in the units of the respective $r_{g,i}/c$ $i$ = $1$, $2$ )."
 Here we allix a gvroradius with the index “gq when it is à value given for the local magnetic field component., Here we affix a gyroradius with the index $g$ ' when it is a value given for the local magnetic field component.
 Index ο means the field including the iclcl perturbations (see below)., Index $e$ ' means the field including the field perturbations (see below).
 The increasing perturbation amplitude is reproduced. in simulations by decreasing the ime period Af between the successive. scatterings., The increasing perturbation amplitude is reproduced in simulations by decreasing the time period $\Delta t$ between the successive scatterings.
 bor simplicitv. we use the same scattering. pattern (AQ andl Af in units of ry fe) upstream and. downstream the shock. eading to the same values of wife) in these regions (sec. however. Ostrowski 1993).," For simplicity, we use the same scattering pattern $\Delta \Omega$ and $\Delta t$ in units of $r_g/c$ ) upstream and downstream the shock, leading to the same values of $\kappa_\perp /
\kappa_\parallel$ in these regions (see, however, Ostrowski 1993)."
" One should note that the xwlicle momentum scattering due to presence of turbulent magnetic field is equivalent to the elfective magnetic field arger than the respective uniform mean component. D, or Bo."," One should note that the particle momentum scattering due to presence of turbulent magnetic field is equivalent to the effective magnetic field larger than the respective uniform mean component, $B_1$ or $B_2$."
 In our model. the ellective field can be estimated as Lt is the lower limit for actual field since the amount of power in perturbations with wave-Iengths smaller than ¢A/ cannot be considered. within such a simple model.," In our model, the effective field can be estimated as It is the lower limit for actual field since the amount of power in perturbations with wave-lengths smaller than $c \, \Delta t$ cannot be considered within such a simple model."
" Below. the derived acceleration time scales are presented in units of the formal clilfusive scale Z5=dla,fU.|eyefl)fe or in units of 7,4/6. in the shock normal rest frame."," Below, the derived acceleration time scales are presented in units of the formal diffusive scale $T_0
\equiv 4( \kappa_{n,1} / U_1 + \kappa_{n,2} / U_2 ) / c$ or in units of $r_{e,1}/c$, in the shock normal rest frame."
 Our numerical calculations involve particles with momenta systematically increasing over several orders of magnitude., Our numerical calculations involve particles with momenta systematically increasing over several orders of magnitude.
 In order to avoid any energv dependent systematic effect. we consider the situation with all spatial and time scales. — defined by the dilfusion. coellicient. the mean time between scatterings and the shock velocity proportional to the particle gvro-radius. ry= ο}. ie. to its momentun.," In order to avoid any energy dependent systematic effect we consider the situation with all spatial and time scales – defined by the diffusion coefficient, the mean time between scatterings and the shock velocity – proportional to the particle gyro-radius, $r_g = p/(eB)$ , i.e. to its momentum."
 For à chosen shock velocity and the magnetic field configuration we inject particles in the shock at some initial, For a chosen shock velocity and the magnetic field configuration we inject particles in the shock at some initial
where A is the intensity or number density of the point process.,where $\lambda$ is the intensity or number density of the point process.
 Subsequent studies using (he Chree-point correlation function frequently used a different. description. of the configuration of triplets. employing (wo distance measures and one angle measure: raris and 8. (he angle between r4» ancl 754.," Subsequent studies using the three-point correlation function frequently used a different description of the configuration of triplets, employing two distance measures and one angle measure: $s=r_{12}, q=r_{23}/r_{12}$ and $\theta$ , the angle between $r_{12}$ and $r_{23}$."
 See. for example. Gazlanagaelal.(2006):IXulkarnietal. (2007).," See, for example, \citet{gaztanaga05,
 nichol06, kulkarni07}."
. Note that in the above notation the angle is subtended at the second point., Note that in the above notation the angle is subtended at the second point.
 We will use the parametrization of (wo distances ancl an angle in this work., We will use the parametrization of two distances and an angle in this work.
 For the study of galaxy. surveys. a related quantity Q. called the reduced. three-point correlation funetion is often used. where The hierarchical form of (3)) was proposed in Peebles&Groth(1975). based on (heir analvses of the Lickand Zwicky catalogs.," For the study of galaxy surveys, a related quantity $Q$, called the reduced three-point correlation function is often used, where The hierarchical form of \ref{eqn:Q}) ) was proposed in \citet{peebles75} based on their analyses of the Lickand Zwicky catalogs."
 It is an empirical form without theoretical support (Jing&Borner1993).. but has been found to hold in other studies citetszapudiül..," It is an empirical form without theoretical support \citep{jing98}, but has been found to hold in other studies \\citet{szapudi01}."
 In analvses of the 21FGRS and SDSS galaxy survevs. there appears to be variation of ( with @ (e.g.Gaztanagaοἱal.2005:Nicholet2006).," In analyses of the 2dFGRS and SDSS galaxy surveys, there appears to be variation of $Q$ with $\theta$ \citep[e.g.][]{gaztanaga05, nichol06}."
. In the statistics literature. the quantity 47? in (2)) above is more commonly expressed in terms of a function g: where we have used two distances and an angle for the parameters of g.," In the statistics literature, the quantity $dP$ in \ref{eqn:threept}) ) above is more commonly expressed in terms of a function $g^{(3)}$: where we have used two distances and an angle for the parameters of $g^{(3)}$."
" Mollerοἱal. refer to A*g"" as the third-product density.", \citet{moller98} refer to $\lambda^3 g^{(3)}$ as the third-product density.
 Molleretal.(1998) also designed a (hirc-order statistic to distinguish. between certain classes of point process models. while llanisch(1983) used a third-order statistic to examine inner linearities in point patterns.," \citet{moller98} also designed a third-order statistic to distinguish between certain classes of point process models, while \citet{hanisch83} used a third-order statistic to examine inner linearities in point patterns."
 See also Sehladitz&Baddeley(2000)., See also \citet{schladitz00}.
.. These third-order statistics are integrated versions of the third-product density. ancl can beconsidered third-order versions of the secouc-orcler A function.," These third-order statistics are integrated versions of the third-product density, and can beconsidered third-order versions of the second-order $K$ function."
" For our purposes. we define such a third-order function. which we denote by A: where (α.ϱ) denotes an interval that includes ὁ but not α and ο=[a,.03].0 dós a range of angles."," For our purposes, we define such a third-order function, which we denote by $\mathcal{K}$: where $(a,b]$ denotes an interval that includes $b$ but not $a$ and $\Omega=[\alpha_1, \alpha_2], 0\le \alpha_1 < \alpha_2 \le \pi$ is a range of angles."
 Thethird-order statistic of Molleretal.(1998). corresponds to AK with A»=0.R4Ry and O= [0.7].," Thethird-order statistic of \citet{moller98} corresponds to $\mathcal{K}$ with $R_2=0, R_3=R_1$ and $\Omega=[0,\pi]$ ."
 The quantity AX(0.Ry}.(ie.ia2H3].Q) has an intuitive ggiven a randomly chosen object at x. APA.H4].(Hs.Rs}.Q) is the expected number of object pairs at y and z such that |y—x|€(0.24].|z(Hs.Rs) ," The quantity $\mathcal{K}((0,R_1],
(R_2,R_3], \Omega)$ has an intuitive given a randomly chosen object at $\mathbf{x}$ , $\lambda^2\mathcal{K}((0,R_1],
(R_2,R_3], \Omega)$ is the expected number of object pairs at $\mathbf{y}$ and $\mathbf{z}$ such that $|\mathbf{y}-\mathbf{x}|\in (0, R_1],
|\mathbf{z}-\mathbf{x}| \in (R_2, R_3]$ "
"an enhancement factor of about 10 in nitrogen in the symbiotic nebulosity from optical, NIR and UV spectroscopy.","an enhancement factor of about 10 in nitrogen in the symbiotic nebulosity from optical, NIR and UV spectroscopy."
" The mean abundances found in M giants show an over-abundance of N by a factor of 3.5 compared to solar values (Smith&Lambert1985,1986),, due to the conversion of carbon into nitrogen and mixing into the atmosphere of the red giant (Nussbaumeretal. 1988)."," The mean abundances found in M giants show an over-abundance of N by a factor of 3.5 compared to solar values \citep{SmL85,SmL86}, due to the conversion of carbon into nitrogen and mixing into the atmosphere of the red giant \citep{NSS88}."
. Larger values may be attributed to more advanced nuclear burning stages., Larger values may be attributed to more advanced nuclear burning stages.
" Therefore the shocked material in the softer component was probably initially part of the red giant wind, which has not been further reprocessed during thermonuclear burning close to the white dwarf (Schmid&Schild1990) and has been ejected outside of the symbiotic nebula within jet."," Therefore the shocked material in the softer component was probably initially part of the red giant wind, which has not been further reprocessed during thermonuclear burning close to the white dwarf \citep{ScS90} and has been ejected outside of the symbiotic nebula within a jet."
 The possible lack of short term variability in athe X-ray light curve may support the above scenario., The possible lack of short term variability in the X-ray light curve may support the above scenario.
" Variability with time scales of minutes to hours is typically observed in symbiotics where X-ray originate in the accretion flow (e.g.SS7317,Ezeetal.2010,orCHCyg,Mukaietal. 2007)."," Variability with time scales of minutes to hours is typically observed in symbiotics where X-ray originate in the accretion flow \citep[e.g. SS73 17, Eze et al. 2010, or CH Cyg,][]{MIK07}."
". However, because the X-ray count rate from CCyg leads to large error bars, we cannot definitely exclude flickering."," However, because the X-ray count rate from Cyg leads to large error bars, we cannot definitely exclude flickering."
 Variability is statistically detected only in the UVW1 band at 2910 and the rms amplitude of this variability is only1%., Variability is statistically detected only in the UVW1 band at 2910 and the rms amplitude of this variability is only.
. We report the detection for the first time of X-ray emission from the symbiotic star V1329 Cyg which was not previously detected byROSAT due to its low X-ray flux., We report the detection for the first time of X-ray emission from the symbiotic star V1329 Cyg which was not previously detected by due to its low X-ray flux.
 The observed X-ray temperatures and especially the unabsorbed nature of the soft component of the X-ray spectrum indicate that some of the X-ray emission might originate in shocks inside a jet outside the symbiotic nebula., The observed X-ray temperatures and especially the unabsorbed nature of the soft component of the X-ray spectrum indicate that some of the X-ray emission might originate in shocks inside a jet outside the symbiotic nebula.
" In this regard, V1329 Cyg adds to the small but growing group of jet-driving symbiotics with X-ray emission."," In this regard, V1329 Cyg adds to the small but growing group of jet-driving symbiotics with X-ray emission."
 MS wishes to thank Hans Martin Schmid for fruitful discussions., MS wishes to thank Hans Martin Schmid for fruitful discussions.
" This work is based on observations obtained withXMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA)."," This work is based on observations obtained with, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA)."
" We thank NASA for funding this work through XMM-Newton AO-8 awards NNXO9AP88G to GJML and JLS, and NNX10AK31G to JLS."," We thank NASA for funding this work through XMM-Newton AO-8 awards NNX09AP88G to GJML and JLS, and NNX10AK31G to JLS."
 We also acknowledge helpful comments and suggestions by an anonymous referee., We also acknowledge helpful comments and suggestions by an anonymous referee.
"microwave background (CMB) (2) is given by Here, the spin temperature, 7;, CMB brightness temperature, T, baryonic overdensity, δε, (which we assume to be equal to the dark matter overdensity) and neutral fraction, zg;, are calculated for a location a at redshift z.","microwave background (CMB) \citep[][]{madau1997} is given by Here, the spin temperature, $T_{\rm s}$, CMB brightness temperature, $T_{\gamma}$ , baryonic overdensity, $\delta_b$, (which we assume to be equal to the dark matter overdensity) and neutral fraction, $x_{H\,{\scriptsize I}\,\,}$, are calculated for a location $\xvec$ at redshift $z$."
" We assume 7;>>T, during the epoch of reionisation (??) and ignore the enhancement of brightness temperature fluctuations due to peculiar velocities in overdense regions (??).."," We assume $T_{\rm s} \gg T_{\gamma}$ during the epoch of reionisation \citep{ciardi2003,furl2006} and ignore the enhancement of brightness temperature fluctuations due to peculiar velocities in overdense regions \citep{bharadwaj2005,bl2005}. ."
" Peculiar velocities were included in the semi-numerical model of ?,, who found their effect to be small on scales 10 Mpc."," Peculiar velocities were included in the semi-numerical model of \cite{mesinger2007}, who found their effect to be small on scales $\sim 10$ Mpc."
 Figure 1 shows an example of a synthetic three-dimensional epoch of reionisation signal data cube at z=7 with and without the effect of limited instrumental angular resolution (see Section 4.2 for a description of the synthesised beam resolution)., Figure \ref{eor_model} shows an example of a synthetic three-dimensional epoch of reionisation signal data cube at $z = 7$ with and without the effect of limited instrumental angular resolution (see Section \ref{The synthesised beam and beam depolarisation} for a description of the synthesised beam resolution).
" For the purpose of demonstrating foreground removal, we do not consider the evolution of the IGM in the line-of-sight direction within the observed bandpass."," For the purpose of demonstrating foreground removal, we do not consider the evolution of the IGM in the line-of-sight direction within the observed bandpass."
" Therefore, the realisation of the density and ionisation state of the IGM in each frequency channel is considered to be at the same stage of cosmic evolution."," Therefore, the realisation of the density and ionisation state of the IGM in each frequency channel is considered to be at the same stage of cosmic evolution."
 This is a valid approximation when the range in redshift corresponding to the simulated bandwidth is small compared with the central redshift of the bandwidth Az/z<1 for redshifts greater than the redshift of overlap £>Zov., This is a valid approximation when the range in redshift corresponding to the simulated bandwidth is small compared with the central redshift of the bandwidth $\Delta z/z \ll 1$ for redshifts greater than the redshift of overlap $z > z_{\rm ov}$.
 We also assume that the reionisation signal is polarised and is therefore present in Stokes I only., We also assume that the reionisation signal is non-polarised and is therefore present in Stokes $I$ only.
 Diffuse Galactic synchrotron emission originates from relativistic electrons in the interstellar medium interacting with the Galactic magnetic field., Diffuse Galactic synchrotron emission originates from relativistic electrons in the interstellar medium interacting with the Galactic magnetic field.
" Although we are primarily concerned with the polarised component in this paper, we discuss our model for total intensity of synchrotron emission here because of the proportionality between the two, as well as to demonstrate the combined effect of continuum and polarised-component foreground removal."," Although we are primarily concerned with the polarised component in this paper, we discuss our model for total intensity of synchrotron emission here because of the proportionality between the two, as well as to demonstrate the combined effect of continuum and polarised-component foreground removal."
" For more detailed non-polarised and polarised foreground models, including synchrotron emission from discrete sources such as supernova remnants and free-free emission from diffuse ionised gas, see ??? Foreground contamination and its removal could have significant consequences for the detectability of the weak, redshifted 21-cm signal."," For more detailed non-polarised and polarised foreground models, including synchrotron emission from discrete sources such as supernova remnants and free-free emission from diffuse ionised gas, see \cite{jelic2008,bowman2009,jelic2010}
 Foreground contamination and its removal could have significant consequences for the detectability of the weak, redshifted 21-cm signal."
" Indeed, the total intensity of foregrounds will be brighter than the cosmological 21-cm signal by 4-5 orders of magnitude."," Indeed, the total intensity of foregrounds will be brighter than the cosmological 21-cm signal by 4–5 orders of magnitude."
" The three main sources of foreground contamination of the 21-cm signal are DGSE (which comprises ~ 70 per cent near 150 MHz), extragalactic point sources (~ 27 per cent) and Galacticbremsstrahlung (~ 1 per cent) (?).."," The three main sources of foreground contamination of the 21-cm signal are DGSE (which comprises $\sim$ 70 per cent near 150 MHz), extragalactic point sources $\sim$ 27 per cent) and Galacticbremsstrahlung $\sim$ 1 per cent) \citep{shaver1999}."
 The frequency dependence of each of these foregrounds can be approximated by a power law with a running spectral index (??)..," The frequency dependence of each of these foregrounds can be approximated by a power law with a running spectral index \citep{shaver1999,tegmark2000}."
" While the sum of power laws is not in general a power law, over a relatively narrow frequency range, B, (such as that considered in this paper where B/v< 1), a Taylor expansion around a power law can be used to describe the spectral shape."," While the sum of power laws is not in general a power law, over a relatively narrow frequency range, $B$, (such as that considered in this paper where $B/\nu\ll 1$ ), a Taylor expansion around a power law can be used to describe the spectral shape."
" We therefore also approximate the sum of foregrounds as a power law with a running spectral index, and specialise to the case of Galactic synchrotron emission, which dominates the foregrounds."," We therefore also approximate the sum of foregrounds as a power law with a running spectral index, and specialise to the case of Galactic synchrotron emission, which dominates the foregrounds."
" For the purpose of demonstrating the removal of polarised foreground leakage, we assume that the brightest point sources have already been removed."," For the purpose of demonstrating the removal of polarised foreground leakage, we assume that the brightest point sources have already been removed."
 The total intensity of Galactic synchrotron emission varies as a function of both sky position and frequency., The total intensity of Galactic synchrotron emission varies as a function of both sky position and frequency.
" We model thefrequency and angular dependence of Galactic synchrotron foreground emission (following??7?) by first constructing a realisation of the angular fluctuations in theforeground, dT (0). at a particular frequency within"," We model thefrequency and angular dependence of Galactic synchrotron foreground emission \citep[following][]{shaver1999,tegmark2000,wang2006} by first constructing a realisation of the angular fluctuations in theforeground, $\delta T^{\rm G}_{\rm b,0}(\vec{\theta})$ , at a particular frequency within"
The last point is interesting. and is highlighted in Figure 5 which also provides a test of our code for general orbits against the semi-analytic case of a conical bow shock.,"The last point is interesting, and is highlighted in Figure \ref{fig5}
 which also provides a test of our code for general orbits against the semi-analytic case of a conical bow shock."
 The figure is organized in four columns., The figure is organized in four columns.
 From bottom to top. the viewing inclination ranges from pole-on to edge-on through angles of 7=0.2.10.30.45.G0.SO.SS. and 90.," From bottom to top, the viewing inclination ranges from pole-on to edge-on through angles of $i=0, 2, 10, 30, 45, 60,
80, 88,$ and $90^\circ$."
 From left to right for the first three columns. the orbital period is | year. 100 years. and 10.000 years.," From left to right for the first three columns, the orbital period is 1 year, 100 years, and 10,000 years."
 Profiles in the far right column are for the purely conical CWIR described in the previous section., Profiles in the far right column are for the purely conical CWIR described in the previous section.
 The orbits are taken as cireular. with roan=0.00100 so that the CWIR exists essentially everywhere.," The orbits are taken as circular, with $r_{\rm orb} = 0.001r_{\rm c}$ so that the CWIR exists essentially everywhere."
 The profile shapes have not been gaussian smoothed., The profile shapes have not been gaussian smoothed.
" The bow shock has .=40° and the compressed layer has 3”45"" in each instance.", The bow shock has $\beta = 40^\circ$ and the compressed layer has $\beta '= 45^\circ$ in each instance.
 Note that the vertical scales (not shown) are the same within a row but vary between rows., Note that the vertical scales (not shown) are the same within a row but vary between rows.
 For the longest period case. the profiles in the third column match those of the strict conical bow shock ease in the fourth column.," For the longest period case, the profiles in the third column match those of the strict conical bow shock case in the fourth column."
 What is surprising is that models with /?=100 years for which the wrapping radius is larger than the eritical radius by two orders of magnitude deviate from the expectations of the conical bow shock ease., What is surprising is that models with $P=100$ years for which the wrapping radius is larger than the critical radius by two orders of magnitude deviate from the expectations of the conical bow shock case.
 The reason is that the emission in the compressed zone is very large., The reason is that the emission in the compressed zone is very large.
 The factor of 4 increase in density combined with the density square emissivity dependence in the outer wind along with the narrowness of the zone leads to relatively “spikey” features in velocity that appear in the line profile., The factor of 4 increase in density combined with the density square emissivity dependence in the outer wind along with the narrowness of the zone leads to relatively `spikey' features in velocity that appear in the line profile.
 Of course with instrumental smearing. the significance of these features will be," Of course with instrumental smearing, the significance of these features will be"
[uin]: the oscillations skew strongly in the wave propagation direction.,", the oscillations skew strongly in the wave propagation direction."
 The dotted. lines correspond. to a superluminal wave Jy=2., The dotted lines correspond to a superluminal wave $\beta_V=2$.
 As UV increases the svstem evolves toward svmunmetric oscillations as shown in the subligure on the right., As $\beta_V$ increases the system evolves toward symmetric oscillations as shown in the subfigure on the right.
 A nonzero initial velocity (Ju=0) also shifts the oscillations Forward: (clash-clottecl lines), A nonzero initial velocity $\beta_{0\pm}>0$ ) also shifts the oscillations forward (dash-dotted lines).
 lt should. be emphasized that we intentionally choose a moderate n/— here to illustrate that an oscillatory solution can exist even for a nmiocerate pair density., It should be emphasized that we intentionally choose a moderate $n_-$ here to illustrate that an oscillatory solution can exist even for a moderate pair density.
 ‘To seek an oscillatorv solution one may define the oscillation periodicity. (in units of ων) where the subscripts a. 6 label respectively the phases at which 3=Guin and o= yas. respectively.," To seek an oscillatory solution one may define the oscillation periodicity (in units of $1/\omega_{_{GJ}}$ ) where the subscripts $a$, $b$ label respectively the phases at which $\beta_-=\beta_{\rm min}$ and $\beta_-=\beta_{\rm max}$ , respectively."
" The oscillation frequency. is then given by w=πω,1. where uuum(L5sLotsNP""EH"," The oscillation frequency is then given by $\omega=2\pi\omega_{_{GJ}}/\hat{T}$, where $\omega_{_{GJ}}\approx (1.5\times10^{11}\,{\rm s}^{-1})P^{-1}_{0.1}B_8$."
 Using2. the approximations. (27))an analvtical solution for a LAEW can be derived., Using the approximations \ref{eq:g2}) )an analytical solution for a LAEW can be derived.
 For Jpc lousing >(€)|g(£)~2|£|/v. one obtains the Following approximation to (26)): where the upper and lower signs correspond to £=£i and £«£1. respectively.," For $\beta_V\gg1$, using $\gamma_-(\xi)+g(\xi)\sim 2|\xi|/\beta_V$, one obtains the following approximation to \ref{eq:Phi2}) ): where the upper and lower signs correspond to $\xi\geq\xi_1$ and $\xi<\xi_1$, respectively."
 For 4 lone max expand. (31)) in L/2y and substitute it for (30)). obtaining A second approximation. applies. for⋅ £4o/2n/—NX»504. in. which case one estimates the frequency as .l2 . ∖∖⋎↓↕⋖⋅↓⋅∢⊾∣⇜≮↙⋎↗∣∶↿∖−≽⊔⊐∣⇜∩↙∣∣↓⊳∖∣⇂↥⋖⊾↓≻↓⋜↧⊳∖⊔↓⋜↧⇂↓⋅⋖⋅⊏↥⋯⋅⊔≼⇍∙∖⇁∪⇂ ⋅ the pair. plasma.," For $\beta_V\gg1$ one may expand \ref{eq:Phi3}) ) in $1/\beta_V$ and substitute it for \ref{eq:periodicity}) ), obtaining A second approximation applies for $\tilde{E}^2_0/2n_-\gg\gamma_{0\pm}$, in which case one estimates the frequency as where $\omega_p=(2n_-)^{1/2}\omega_{_{GJ}}$ is the plasma frequency of the pair plasma."
" As an example. one has w=10""Ms for n.=5IO. μας=40"" and wes,=10s+."," As an example, one has $\omega\approx 10^9\,{\rm s}^{-1}$ for $n_-=10^2$, $u_{\rm max}=10^6$ and $\omega_{GJ}=10^{11}\,{\rm s}^{-1}$."
 The oscillation frequency decreases às μμ dnereases. which can be understood. as an increase in the elective mass of electrons or positrons.," The oscillation frequency decreases as $u_{\rm max}$ increases, which can be understood as an increase in the effective mass of electrons or positrons."
" Vhe condition. for 4 Ods|£]<&m£,3Mbfan this gives tun,mo£5/4nοMyine", The condition for $\Phi>0$ is $|\xi|<\xi_b\approx-\xi_a\approx\beta_V\tilde{E}^2_0/4n_-$ ; this gives $u_{\rm max}\approx \tilde{E}^2_0/4n_-\approx-u_{\rm min}$.
 Note that such symmetry in oscillation is a direct. consequence of our assumption of large y and i;=0., Note that such symmetry in oscillation is a direct consequence of our assumption of large $\beta_V$ and $\beta_{\pm0}=0$.
" Phe solution for QEON«NymJvT ds found to be where d is given by (31)). Oy=9(1). 9,—o(£,0)."," The solution for $0\leq \chi<\chi_{_T}=\beta_V\hat{T}$ is found to be where $\Phi$ is given by \ref{eq:Phi3}) ), $\Phi_0=\Phi(1)$, $\Phi_{a,b}=\Phi(\xi_{a,b})$."
 The electric Bfeld is obtained as The electric field displays a sawtooth wave form. which can be understood. qualitatively in terms of the extreme relativistic limit.," The electric field is obtained as The electric field displays a sawtooth wave form, which can be understood qualitatively in terms of the extreme relativistic limit."
 In this limit positrons and electrons. are accelerated. in opposite directions. giving rise to a current Uil~2)3[n[joi]1.," In this limit positrons and electrons are accelerated in opposite directions, giving rise to a current $|j_\parallel|\sim 2|\beta_-|n_-\gg |j_{0\parallel}|\sim 1$."
 Thus. the electric eld is Lyoxce~Eenxfj with Jo|~1. which reproduces the sawtooth wave form given by (35)).," Thus, the electric field is $E_\parallel\propto \pm |j_\parallel|/\beta_V\sim\pm 2n_-\chi/\beta_V$ with $|\beta_-|\sim1$, which reproduces the sawtooth wave form given by \ref{eq:E}) )."
 An example of anumerical integration of (22)) and (23)) is shown in figure 5. for pairs with an initial. Forward. velocity.," An example of anumerical integration of \ref{eq:EqMotion4}) ) and \ref{eq:waveE4}) ) is shown in figure \ref{fig:euchi} for pairs with an initial, forward velocity."
 Figure ο shows oscillations for particles with an initial velocity toward the star., Figure \ref{fig:uchi} shows oscillations for particles with an initial velocity toward the star.
 Although our analytical solution is obtained [or n1. here we again choose a moderate n in the numerical calculation to show that our oscillatory solution is also valid for n4~1.," Although our analytical solution is obtained for $n_-\gg1$, here we again choose a moderate $n_-$ in the numerical calculation to show that our oscillatory solution is also valid for $n_\pm\sim1$."
 The wave form is similar to that predicted from the numerical moclel (Levinsonctal.2005)., The wave form is similar to that predicted from the numerical model \citep{letal05}.
. The characteristics of the oscillations. including the periodicity ancl amplitude. are independent of the sign of Ly and are not sensitive to the initial conditions τος provided that ros]Mun.," The characteristics of the oscillations, including the periodicity and amplitude, are independent of the sign of $\tilde{E}_0$ and are not sensitive to the initial conditions $\beta_{0\pm}$ provided that $|u_{0\pm}|\ll u_{\rm
max}$ ."
 As no radiative loss. pair production nor wave damping is included. the wave amplitude remains constant.," As no radiative loss, pair production nor wave damping is included, the wave amplitude remains constant."
 In the figure we assume an initial electric field much lower than the typical vacuum. field Eau~3 where Boz4410T.," In the figure we assume an initial electric field much lower than the typical vacuum field $\tilde{E}_{\rm max}\sim 3\times10^6(B/B_c)^{1/2}
P^{-1/2}_{0.1}$ , where $B_c\approx 4.4\times10^9\, \rm T$."
 In. practice. the initial field should benear the pair creation threshold.," In practice, the initial field should benear the pair creation threshold."
" Assuming the pair production threshold to be 5,,. one has Ly8oLonπλτσ/10"")7. (sce discusssion in Sec."," Assuming the pair production threshold to be $\gamma_{th}$, one has $\tilde{E}_0\approx3\times10^4(n_-/5)^{1/2}(\gamma_{th}/10^6)^{1/2}$ (see discusssion in Sec."
 5)., 5).
 Since |ui9[rus]: the oscillating particles have a net Forward [ow velocity.," Since $|u_{\rm min}|<|u_{\rm max}|$, the oscillating particles have a net forward flow velocity."
 The drift momentum can be obtained by averaging vw over one period. (2): where we expand ὁ on 1l and change the integration variable to dy=-, The drift momentum can be obtained by averaging $u_-$ over one period $\hat{T}$ ): where we expand $\Phi$ on $1/\beta_V\ll1$ and change the integration variable to $d\chi=-d\xi/\Phi^{1/2}(\xi)$.
" Terms of order lio)|e] and [n,;,]/n41 are ignored.", Terms of order $|j_{0\parallel}|/n_\pm\ll1$ and $|\eta_{_{GJ}}|/n_\pm\ll1$ are ignored.
 Then both electrons ancl positrons are cragged along in the wave at the same drift velocity p=1(liu)DAL/D, Then both electrons and positrons are dragged along in the wave at the same drift velocity $\beta_D=\bar{u}/(1+\bar{u}^2)^{1/2}$.
 qx 7.LborE5/2nEAMX- one has The drift velocity. decreases as the wave phase speed increases., For $\tilde{E}^2_0/2n_-\gg\gamma_{\pm0}$ one has The drift velocity decreases as the wave phase speed increases.
 In the limit ;» oscillations are purely temporal., In the limit $\beta_V\to\infty$ oscillations are purely temporal.
" The upper anc lower limits to the particle's momentum ave mua=(Plowio|""o)/2 and ""uu=(Pony wog)f2.", The upper and lower limits to the particle's momentum are $u_{\rm max}=(\Gamma+u_{+0}+u_{-0})/2$ and $u_{\rm min}=-(\Gamma-u_{+0}-u_{-0})/2$ .
 The drift. velocity (36)). reduces to ues(noa|owuq)/2., The drift velocity \ref{eq:ubar}) ) reduces to $\bar{u}=(u_{+0}+u_{-0})/2$.
 For we»0-0. particles oscillate svmnmetrically. between tain©Minas0]FSpesο and mua," For $u_{\pm0}=0$, particles oscillate symmetrically between $u_{\rm min}\approx-u_{\rm max}\approx \tilde{E}^2_0/4n_-$ and $u_{\rm max}$."
 ]t is interesting to note that the proportionality L1/: in (37)) is similar to that. predicted. from the single-particletreatment (Rowe 1992b).., It is interesting to note that the proportionality $1/\beta_V$ in \ref{eq:ubar2}) ) is similar to that predicted from the single-particletreatment \citep{r92b}. .
 Llowever. since the single-particle formalism does not include the feedback. ellect of particles on the wave. it predicts a low drift velocity 2p2 L/29.," However, since the single-particle formalism does not include the feedback effect of particles on the wave, it predicts a low drift velocity $\beta_D\approx1/\beta_V$ ."
 Our extact treament shows that thedrift motion can be highly relativisticwith sy—1/(1IVS15)πε]9 , Our extact treament shows that thedrift motion can be highly relativisticwith $\gamma_D\equiv1/(1-\beta^2_D)^{1/2}\approx |\bar{u}|\gg1$ .
LAs a result. the LAEW can drive a relativistic outflow of particles even when the particles are initially at rest.," As a result, the LAEW can drive a relativistic outflow of particles even when the particles are initially at rest."
 Thus. the oscillating gap can supply relativistic pairs to the pulsar wind.," Thus, the oscillating gap can supply relativistic pairs to the pulsar wind."
At the time of this writing several large projects are being pursued in order to clirectly detect! astrophysical eravitational waves.,At the time of this writing several large projects are being pursued in order to directly detect astrophysical gravitational waves.
 This paper concerns a program to detect gravitational waves using pulsars as nearlv-perfect Einstein clocks., This paper concerns a program to detect gravitational waves using pulsars as nearly-perfect Einstein clocks.
 The practical idea is to time a set of milliscconcl pulsars (called the “Pulsar Timing Array”. or PPA) over a number of vears (7)..," The practical idea is to time a set of millisecond pulsars (called the “Pulsar Timing Array”, or PTA) over a number of years \citep{Foster}."
 Some of the millisecond pulsars create pulse trains of exceptional regularity., Some of the millisecond pulsars create pulse trains of exceptional regularity.
 By perturbing the space-time between a pulsar ancl the Earth. the gravitational waves (GWs) will cause extra deviations from the periodicitv in the pulse arrival times (?77)..," By perturbing the space-time between a pulsar and the Earth, the gravitational waves (GWs) will cause extra deviations from the periodicity in the pulse arrival times \citep{Estabrook,
 Sazhin, Detweiler}."
" Thus from the measurements of these deviations (called ""timing-residuals or Tli). one may measure the eravitational waves."," Thus from the measurements of these deviations (called “timing-residuals”, or TR), one may measure the gravitational waves."
 Currently.. several PTA project are operating around the globe.," Currently, several PTA project are operating around the globe."
 Firstly. at the Arecibo Itadio Telescope in North-America several millisecond pulsars have been timed for a number of vears.," Firstly, at the Arecibo Radio Telescope in North-America several millisecond pulsars have been timed for a number of years."
 These observations have already been used to place interesting upper limits on the intensity of gravitational waves which are passing through the Galaxy (??)..," These observations have already been used to place interesting upper limits on the intensity of gravitational waves which are passing through the Galaxy \citep{Kaspi, Lommen}."
 Together with the Green Dank Telescope. the Arecibo Radio Telescope will be used. as an instrument of NANOCrav. the North American PPA.," Together with the Green Bank Telescope, the Arecibo Radio Telescope will be used as an instrument of NANOGrav, the North American PTA."
 Secondly. the European PPA is being set up as an international collaboration between Great. Britain. France. Netherlands. Germany. and. Italy. and. will use 5 European radio telescopes to monitor about 20 milliseconcl pulsars (?)..," Secondly, the European PTA is being set up as an international collaboration between Great Britain, France, Netherlands, Germany, and Italy, and will use 5 European radio telescopes to monitor about 20 millisecond pulsars \citep{Stappers}."
 Finally. the Parkes PTA in Australia has been using the Parkes multi-beam racio-telescope to monitor 20 millisecond pulsars (2)..," Finally, the Parkes PTA in Australia has been using the Parkes multi-beam radio-telescope to monitor 20 millisecond pulsars \citep{Manchester}."
 Some of the Parkes ancl Arecibo data have also been used to place the most stringent limits on the GAB to date (?).., Some of the Parkes and Arecibo data have also been used to place the most stringent limits on the GWB to date \citep{Jenet-2006}.
 One of the main astrophysical targets of the Ας is the stochastic background of the gravitational waves (GWD)., One of the main astrophysical targets of the PTAs is the stochastic background of the gravitational waves (GWB).
 This GWD is thought to be generated. by a [arge. number of black-hole binaries which are thought to be located. at, This GWB is thought to be generated by a large number of black-hole binaries which are thought to be located at
“gap” in the halo MDF with between —4.0 and —5.0.,“gap” in the halo MDF with between $-4.0$ and $-5.0$.
 It adoptes a scenario of negative feedback from Population III stars., It adoptes a scenario of negative feedback from Population III stars.
 Figure 10 suggests that it only roughly fits the portion of the MDF with—3., Figure \ref{HES_Ks} suggests that it only roughly fits the portion of the MDF with.
0.. It also fails to predict the sharp drop at the low-metallicity end as well., It also fails to predict the sharp drop at the low-metallicity end as well.
 Another model that has been tested is GAlaxy MErger Tree and Evolution (GAMETE—?).., 	 Another model that has been tested is GAlaxy MErger Tree and Evolution \citep[GAMETE -- ][]{Salvadori2007MNRAS}.
 It is a Monte Carlo code to reconstruct the merger tree of the Milky Way and to follow the evolution of gas and stars along the tree., It is a Monte Carlo code to reconstruct the merger tree of the Milky Way and to follow the evolution of gas and stars along the tree.
" This model defines an input parameter, the critical metallicity Z,,, which governs the transition from Pop III to Pop II star formation."," This model defines an input parameter, the critical metallicity $Z_{cr}$, which governs the transition from Pop III to Pop II star formation."
" We compare our observed MDF with the simulated results corresponding to different values of Z,,, as shown in Figure 11.."," We compare our observed MDF with the simulated results corresponding to different values of $Z_{cr}$, as shown in Figure \ref{HES_Zcr}."
" Although according to the observational data available at that time, Z,,=107Z; was regarded as the fiducial model, it obviously cannot fit our observations here."," Although according to the observational data available at that time, $Z_{cr}=10^{-4}Z_{\odot}$ was regarded as the fiducial model, it obviously cannot fit our observations here."
 All the predictions fail to fit the location of the peak of the observed MDF., All the predictions fail to fit the location of the peak of the observed MDF.
" Similarly to the conclusions in Paper V, the model with Z,,=10774Z, appears to partially fit our observed MDF, being able to reproduce the tail with and best predict the sharp drop at--2."," Similarly to the conclusions in Paper V, the model with $Z_{cr}=10^{-3.4}Z_{\odot}$ appears to partially fit our observed MDF, being able to reproduce the tail with and best predict the sharp drop at."
"6.. Based on the (for now) largest metal-poor turnoff-star sample from the HES database and moderate-resolution follow-up observations, we have statistically investigated the MDF of local MSTO stars in the Galactic halo."," 	 Based on the (for now) largest metal-poor turnoff-star sample from the HES database and moderate-resolution follow-up observations, we have statistically investigated the MDF of local MSTO stars in the Galactic halo."
" Considering the fact that our sample mainly consists of unevolved main-sequence (and subgiant) stars with low metallicities, it could also provide additional useful information on Galactic chemical evolution."," 	 Considering the fact that our sample mainly consists of unevolved main-sequence (and subgiant) stars with low metallicities, it could also provide additional useful information on Galactic chemical evolution."
" For example, a kinematic analysis of this sample could be used to re-visit the role of accretion of the interstellar medium during the long lifetimes of metal-poor stars, as approximately calculated in a number of early works (e.g., ???)) and also discussed by more recent studies (e.g, ???))."," For example, a kinematic analysis of this sample could be used to re-visit the role of accretion of the interstellar medium during the long lifetimes of metal-poor stars, as approximately calculated in a number of early works (e.g., \citealt{Talbot1977ApJS,Yoshii1981AA,Iben1983MmSAI}) ) and also discussed by more recent studies (e.g, \citealt{Christlieb2004ApJ,Norris2007ApJ,Frebel2009MNRAS}) )."
" It should also be pointed out that all of our comparisons of the MDFs have been performed under the assumption that we are modeling a halo population, which current evidence suggests is an over-simplification."," It should also be pointed out that all of our comparisons of the MDFs have been performed under the assumption that we are modeling a halo population, which current evidence suggests is an over-simplification."
 It seems likely that the observed MDFs for both the HES MSTO stars and, It seems likely that the observed MDFs for both the HES MSTO stars and
"Stellar kinematics show that there are =2.6«109AF, within z0.015 pe of the Calactic Center (Eckart Cenzel 1997. Chez et al.","Stellar kinematics show that there are $\approx 2.6 \times 10^6
M_\odot$ within $\approx 0.015$ pc of the Galactic Center (Eckart Genzel 1997, Ghez et al."
 1998). centered on the radio source Ser A* (Moeuten et al.," 1998), centered on the radio source Sgr A* (Menten et al."
 1997)., 1997).
 The most plausible explanation is that Ser is a z2.6<LOCAL. accreting black hole., The most plausible explanation is that Sgr $^*$ is a $\approx 2.6 \times 10^6 M_\odot$ accreting black hole.
 Ser is believed to accrete the winds from nearby (~U.l pe} iassive stars (Arabbe ct al., Sgr $^*$ is believed to accrete the winds from nearby $\sim 0.1$ pc) massive stars (Krabbe et al.
 1991)., 1991).
" Uvdrodvnamical simulations of the Galactic Center region. asstuning 10 randomly distributed point sources (to nimic he nearby stars) aud a 10937. black hole. yield accretion rates c124.101A,we| (Coker Melia 1997: jereafter CAD."," Hydrodynamical simulations of the Galactic Center region, assuming 10 randomly distributed point sources (to mimic the nearby stars) and a $10^6 M_\odot$ black hole, yield accretion rates $\approx 1-2 \times 10^{-4} \mpy$ (Coker Melia 1997; hereafter CM)."
 Tf we naively increase this bv a (Doudi capture) factor of 2.6? to account for the actual πας of the central object. the theoretically predicted accretion rate jocomies ~LOOAL.ve3 (but soe 833).," If we naively increase this by a (Bondi capture) factor of $2.6^2$ to account for the actual mass of the central object, the theoretically predicted accretion rate becomes $\sim
10^{-3} \mpy$ (but see 3)."
 The bolometric huninositv of Ser is observed to ος107 ores s| (Cenzel et al.," The bolometric luminosity of Sgr $^*$ is observed to be $\lsim
10^{37}$ ergs $^{-1}$ (Genzel et al."
 1991. Narayan et al.," 1994, Narayan et al."
 1998)., 1998).
 For an accretion rate of ~10°AL.xv|. this corresponds to a radiative efficiency of ~10.*1," For an accretion rate of $\sim 10^{-3} \mpy$, this corresponds to a radiative efficiency of $\sim 10^{-7}$!"
 A possible explanation for the low hunuinositv of Ser is that the accretion. occurs via an advectiou-cdominated accretion flow (ADAE: Naravan. Yi Mahadevan 1995. Mamumoto et al.," A possible explanation for the low luminosity of Sgr $^*$ is that the accretion occurs via an advection-dominated accretion flow (ADAF; Narayan, Yi Mahadevan 1995, Manmoto et al."
 1997. Naravan ct al.," 1997, Narayan et al."
 1998: sec. e.g.. Melia 1992. 199L. Mastichiadis Ozernov 1991. Faleke 1996. Doeckert Duschl 1997 for alternative models of Ser A).," 1998; see, e.g., Melia 1992, 1994, Mastichiadis Ozernoy 1994, Falcke 1996, Beckert Duschl 1997 for alternative models of Sgr $^*$ )."
 All ADAF models in the literature. however. require au accretion rate of x10CAL.vv| (see Table 2 of Quatacrt Narayvau 1999. which gives accretion rates in Eddington wits: their estimates need to be multiplied by a factor of 0.055 for Αινι Hy.," All ADAF models in the literature, however, require an accretion rate of $\lsim 10^{-5} \mpy$ (see Table 2 of Quataert Narayan 1999, which gives accretion rates in Eddington units; their estimates need to be multiplied by a factor of 0.055 for $\mpy$ )."
 HE CALS black hole mass scaled Boudi capture estimate of ~LOPAL.vy.| is taken at face value. there is a discrepancy of a factor of 2107 3n mass accretion rates.," If CM's black hole mass scaled Bondi capture estimate of $\sim 10^{-3} \mpy$ is taken at face value, there is a discrepancy of a factor of $\gsim 10^2$ in mass accretion rates."
 Iu 822. we sharpen this discrepancy by using spectral observations of the Galactic Center to constrain the accretion rate of gas at large distances from the central black hole.," In 2, we sharpen this discrepancy by using spectral observations of the Galactic Center to constrain the accretion rate of gas at large distances from the central black hole."
 The argument is uearly indepeudent of the accretion inodel employed., The argument is nearly independent of the accretion model employed.
 We then calculate Doucdi capture estimates of accretion onto Ser explicitly using the observed spatial distribution and wind propertics of stars in the Calactic Center (853).," We then calculate Bondi capture estimates of accretion onto Sgr $^*$, explicitly using the observed spatial distribution and wind properties of stars in the Galactic Center 3)."
 In sli 855 we conclude and assess the aerecment between various theoretical estimates of the accretion rate in Sey A”., In 4 5 we conclude and assess the agreement between various theoretical estimates of the accretion rate in Sgr $^*$.
 We use X-ray and IR limits on the cussion from the Galactic Center to derive constraints on the accretion rate of optically thin or thick eas outo Ser A., We use X-ray and IR limits on the emission from the Galactic Center to derive constraints on the accretion rate of optically thin or thick gas onto Sgr $^*$.
 We take the distance to the Galactic Ceuter to be 8.0 kpe., We take the distance to the Galactic Center to be $8.0$ kpc.
" Iu what follows. mas denotes the mass of the black hole im uuits of 2.6«10A7.. R denotes the pliysieal radius iu the aceretion flow. r (=R/T.38«10Haps cm) denotes the radius in Sclawarzschild units. aud ri is the accretion rate in Eddington units. where M,=0.055ingM. 1"," In what follows, $m_{S}$ denotes the mass of the black hole in units of $2.6 \times 10^6 M_\odot$, $R$ denotes the physical radius in the accretion flow, $r$ $=R/7.38\times10^{11}m_S$ cm) denotes the radius in Schwarzschild units, and $\dot m$ is the accretion rate in Eddington units, where $\dot M_{edd} = 0.055 m_{S}
M_\odot$ $^{-1}$."
 If the eas accreting outo Ser is optically thin. it will enüt X-rays bv breisstrallune cussion.," If the gas accreting onto Sgr $^*$ is optically thin, it will emit X-rays by bremsstrahlung emission."
 The observed limits on X-ray cinission from Ser can therefore be used to coustrain the accretion rate., The observed limits on X-ray emission from Sgr $^*$ can therefore be used to constrain the accretion rate.
" The accretion rate of the gas is related to the density. p. by the coutinuity equation AM -- Iz R?pv,p] uu where à,10°. the ""outer radius of the flow. is roughly the Boudi capture radius of the black hole."," The accretion rate of the gas is related to the density, $\rho$, by the continuity equation M = 4 R^2 _H _v = M_o (r/r_o )^p, where $r_o \sim 10^5$, the “outer” radius of the flow, is roughly the Bondi capture radius of the black hole."
" The radial velocity of the gas is related to the free fall velocity. 6,y. by Jelκο cud the vertical scale height of the gas is II=ggR."," The radial velocity of the gas is related to the free fall velocity, $v_{ff}$, by $|v| = \eta_v v_{ff}$ and the vertical scale height of the gas is $H =
\eta_H R$."
 Foy a pure Dondi Sow. gy=yy1. while for an ADAF. jjz6 (the viscosity parameter) aud 45;21/2.," For a pure Bondi flow, $\eta_v = \eta_H = 1$, while for an ADAF, $\eta_v \approx \alpha$ (the viscosity parameter) and $\eta_H \approx 1/2$."
 Equation (1)) allows for the possibility that the mass accretion rate iuav decrease with radius due to an outflow/wind (mn which case the assumed geometry is a “thick” equatorial flow with a bipolar outhowing wind: cf Blandford Beechuan 1999)., Equation \ref{cont}) ) allows for the possibility that the mass accretion rate may decrease with radius due to an outflow/wind (in which case the assumed geometry is a “thick” equatorial inflow with a bipolar outflowing wind; cf Blandford Begelman 1999).
 The accretion rate at the, The accretion rate at the
"component show lower velocity dispersions 1n c7, and ορ than the spheroid model particles.",component show lower velocity dispersions in $\sigma_r$ and $\sigma_b$ than the spheroid model particles.
 As expected the model does not show any correlations in the velocity components of the spheroid particles. in line with our metal poor sample.," As expected the model does not show any correlations in the velocity components of the spheroid particles, in line with our metal poor sample."
 Figure 8 shows the strong decrease in radial velocity dispersion. predicted by the Fux model for the disc/bar population while moving away from the plane., Figure \ref{fig:compFuxSigVr} shows the strong decrease in radial velocity dispersion predicted by the Fux model for the disc/bar population while moving away from the plane.
 In the plane the predicted dispersion ts larger for the disc/bar than the spheroid one. due to the influence of the bar driven orbits.," In the plane the predicted dispersion is larger for the disc/bar than the spheroid one, due to the influence of the bar driven orbits."
" At high latitude the c, dispersion is small for the bar/dise componer= following a dise like behaviour."," At high latitude the $\sigma_r$ dispersion is small for the bar/disc component, following a disc like behaviour."
 The spheroid population keeps the same velocity dispersion along the minor axis., The spheroid population keeps the same velocity dispersion along the minor axis.
 In Fig., In Fig.
 5 we have over-plotted the Fux model radial velocity dispersions with our data for the full sample. the metal rich and the metal poor parts. as defined by the and quantiles in the [Fe/H] distribution.," \ref{fig:compFuxSigVr} we have over-plotted the Fux model radial velocity dispersions with our data for the full sample, the metal rich and the metal poor parts, as defined by the and quantiles in the [Fe/H] distribution."
 For the metal rich part we observe a strong decrease of the metallicity dispersion while going away from the plane. coherent with the behaviour of the dise/bar particles of Fux model.," For the metal rich part we observe a strong decrease of the metallicity dispersion while going away from the plane, coherent with the behaviour of the disc/bar particles of Fux model."
 For the metal poor part we observe a constant velocity dispersion. as for the spheroid particles of the Fux model. but with a mean velocity dispersion of about 20 km/s lower than the model.," For the metal poor part we observe a constant velocity dispersion, as for the spheroid particles of the Fux model, but with a mean velocity dispersion of about 20 km/s lower than the model."
 It is this difference of velocity dispersion for the metal poor component which leads to a too high global velocity dispersion of the Fux model compared to the BRAVA data. leading ? to an interpretation different from ours: at b=—8? the radial velocity distribution is compatible with the disc/bar component of the Fux model without the need for an old spheroid component.," It is this difference of velocity dispersion for the metal poor component which leads to a too high global velocity dispersion of the Fux model compared to the BRAVA data, leading \cite{Howard09} to an interpretation different from ours: at $b=-8\degr$ the radial velocity distribution is compatible with the disc/bar component of the Fux model without the need for an old spheroid component."
 However with a velocity dispersion for the spheroid of ~ 100 km/s instead of ~ 120 km/s the global velocity dispersion would be coherent with the full Fux model (spheroid+disce)., However with a velocity dispersion for the spheroid of $\sim$ 100 km/s instead of $\sim$ 120 km/s the global velocity dispersion would be coherent with the full Fux model (spheroid+disc).
" The use of the Besangoon model adds support to the interpretation that the decrease 1n c, with latitude. seen in disc/bar particles of ? and in our metal-rich samples. can be due to the disc replacing the bar in the sample: according to the Besangoon model. the thin disc contamination of our sample at b=-—6° was expected to be of with ao, = 62 km/s and at b=—12* with c, = 47 km/s. Our analysis of the kinematics as a function of metallicity in Baade's Window shows that our sample can be decomposed in two distinet populations for which we suggest different formation scenaril."," The use of the Besançoon model adds support to the interpretation that the decrease in $\sigma_{r}$ with latitude, seen in disc/bar particles of \cite{Fux99} and in our metal-rich samples, can be due to the disc replacing the bar in the sample: according to the Besançoon model, the thin disc contamination of our sample at $b=-6\degr$ was expected to be of with a $\sigma_{r}$ = 62 km/s and at $b=-12\degr$ with $\sigma_{r}$ = 47 km/s. Our analysis of the kinematics as a function of metallicity in Baade's Window shows that our sample can be decomposed in two distinct populations for which we suggest different formation scenarii."
 The metal poor component does not show any correlation between its velocity components and is therefore consistent with an isotropic rotating population., The metal poor component does not show any correlation between its velocity components and is therefore consistent with an isotropic rotating population.
 perΠΙ showed that this component is enriched in |Mg/Fe]., II showed that this component is enriched in [Mg/Fe].
 We interpret this population as an old spheroid with a rapid time-scale formation., We interpret this population as an old spheroid with a rapid time-scale formation.
 The metal rich component shows a vertex deviation consistent with that expected from tags a population with orbits supporting a bar., The metal rich component shows a vertex deviation consistent with that expected from tags a population with orbits supporting a bar.
 PaperIL showed that this component has a [Mg/Fe] near solar., II showed that this component has a [Mg/Fe] near solar.
 We interpret this population as a pseudo-bulge formed over a long time scale through dise secular evolution under the action of a bar., We interpret this population as a pseudo-bulge formed over a long time scale through disc secular evolution under the action of a bar.
 This pseudo-bulge is gradually disappearing when moving away from the Galactic plane., This pseudo-bulge is gradually disappearing when moving away from the Galactic plane.
 In this context we can give à new consistent interpretation of the metallicity gradient in the bulge., In this context we can give a new consistent interpretation of the metallicity gradient in the bulge.
 A metallicity gradient is indeed visible in the bulge when observing further away from the plane than Baade's Window (? and ID. while in the inner regions (Jo]<4°) no gradient in metallicity has been found (? and ?)).," A metallicity gradient is indeed visible in the bulge when observing further away from the plane than Baade's Window \cite{Frogel99} and I), while in the inner regions $\vert b \vert\leq4\degr$ ) no gradient in metallicity has been found \cite{Ramirez00} and \cite{Rich07}) )."
 This can be understood if the bar as well as the old spheroid population are both present in the inner regions. leading to a constant metallicity at |b]<4°. while the bar influence gradually fades further away from the plane than Baade’s Window.," This can be understood if the bar as well as the old spheroid population are both present in the inner regions, leading to a constant metallicity at $\vert b \vert\leq4\degr$, while the bar influence gradually fades further away from the plane than Baade's Window."
 At high latitudes only the old spheroid remains: the mean metallicity of the outer bulge measured by ? of [Fe/H|~—0.3 dex corresponds very well to our metal poor population., At high latitudes only the old spheroid remains: the mean metallicity of the outer bulge measured by \cite{IbataGilmore95} of $\sim-0.3$ dex corresponds very well to our metal poor population.
 This scenario is also consistent with the distribution of bulge globular clusters of ?. who found no evidence for a metallicity gradient but all their clusters with [Fe/H|»—0.5 dex wre located within ([5]€3°)., This scenario is also consistent with the distribution of bulge globular clusters of \cite{Valenti09} who found no evidence for a metallicity gradient but all their clusters with $>-0.5$ dex are located within $\vert b \vert\leq5\degr$ ).
 In what concerns the age of the two populations. we expect the spheroid component to be old while the pseudo-bulge component may contain both the old stars of the inner dise redistributed by the bar and younger stars whose formation has been triggered by the bar gas flow.," In what concerns the age of the two populations, we expect the spheroid component to be old while the pseudo-bulge component may contain both the old stars of the inner disc redistributed by the bar and younger stars whose formation has been triggered by the bar gas flow."
 ? found that although the bulk of the bulge population is old. a fraction of the stars are of intermediate age (1 to 7 Gyr).," \cite{vanLoon03} found that although the bulk of the bulge population is old, a fraction of the stars are of intermediate age (1 to 7 Gyr)."
 ? observed Mira stars of ages 1-3 Gyr at all latitudes from —1.2 to —5.8 in the OGLE-II data., \cite{GroenewegenBlommaert05} observed Mira stars of ages 1-3 Gyr at all latitudes from $-1.2$ to $-5.8$ in the OGLE-II data.
 ? found 4 bulge stars with ages lower than 3 Gyr at a latitude of b=—107., \cite{Uttenthaler07} found 4 bulge stars with ages lower than 3 Gyr at a latitude of $b=-10\degr$.
 ? found 3 microlensed bulge dwarfs with ages lower than 5 Gyr., \cite{Bensby09} found 3 microlensed bulge dwarfs with ages lower than 5 Gyr.
 of the variable stars detected by ? are distributed within |5|<5° and most of them should be large-amplitude and long-period variables such as Mira variables or OH/IR stars., of the variable stars detected by \cite{KouzumaYamaoka09} are distributed within $\vert b \vert\leq5\degr$ and most of them should be large-amplitude and long-period variables such as Mira variables or OH/IR stars.
 This intermediate age population has been shown to trace the Galactic bar (?.. ?..?...2)). although providing a larger bar angle (~ 40°) than studies based on older tracers such as red clump stars (.20°. ?..?)).," This intermediate age population has been shown to trace the Galactic bar \cite{vanLoon03}, \cite{Izumiura95paperIII}, \cite{GroenewegenBlommaert05}, \cite{KouzumaYamaoka09}) ), although providing a larger bar angle $\sim40\degr$ ) than studies based on older tracers such as red clump stars $\sim20\degr$, \cite{Stanek94}, \cite{Babusiaux05}) )."
 This discrepancy in the bar angle could well be explained if the old tracers probed a mix of spheroid and bar structures while the young tracers only probe the bar one (although biases or the longitude area surveyed needs also to be taken into account. see ?)).," This discrepancy in the bar angle could well be explained if the old tracers probed a mix of spheroid and bar structures while the young tracers only probe the bar one (although biases on the longitude area surveyed needs also to be taken into account, see \cite{Nishiyama05}) )."
 If this intermediate age population was associated to a part of the bar component. their presence in the CMDs would decrease while going away from the plane as the main bar component and would therefore be a small fraction of the CMD of ? at b=-6°.," If this intermediate age population was associated to a part of the bar component, their presence in the CMDs would decrease while going away from the plane as the main bar component and would therefore be a small fraction of the CMD of \cite{Zoccali03} at $b=-6\degr$."
 ? obtained à proper motion decontaminated CMD with a well defined old turn-off in an inner field (/=|. b= —37).," \cite{Clarkson08} obtained a proper motion decontaminated CMD with a well defined old turn-off in an inner field $l=1\degr$, $b=-3\degr$ )."
 However we would expect an intermediate age population associated with the bar to be metal rich. which. due to the age-metallicity degeneracy. would imply that this population could be hidden in the CMD of ? if its contribution is small enough compared to the bulk of the bulge population.," However we would expect an intermediate age population associated with the bar to be metal rich, which, due to the age-metallicity degeneracy, would imply that this population could be hidden in the CMD of \cite{Clarkson08} if its contribution is small enough compared to the bulk of the bulge population."
 The new filter combination proposed by the ACS Bulge Treasury Programme to break the age-netallicity-temperature degeneracy (?) should provide new lights on this issue., The new filter combination proposed by the ACS Bulge Treasury Programme to break the age-metallicity-temperature degeneracy \citep{Brown09} should provide new lights on this issue.
 We note that in Baade's Window. neither the kinematics hor the chemistry allows to distinguish what we call the old spheroid to the thick disc.," We note that in Baade's Window, neither the kinematics nor the chemistry allows to distinguish what we call the old spheroid to the thick disc."
 The mean metallicity of the solar neighbourhood thick dise is however lower (e.g. ? derived Fe/H] =-0.6 and [Mg/H] = -0.2) than the mean metallicity, The mean metallicity of the solar neighbourhood thick disc is however lower (e.g. \cite{Fuhrmann08} derived [Fe/H] $=-0.6$ and [Mg/H] $=-0.2$ ) than the mean metallicity
towards the earth would produce brighter GRBs. and. eireularlv. polarized GWs.,"towards the earth would produce brighter GRBs, and circularly polarized GWs."
" As the viewing angle becomes larger relative to (he polar axis. (he luminosity of (he GRB decreases x@7. while the degree of linear polarization increases as I?~10(0/00degree),"," As the viewing angle becomes larger relative to the polar axis, the luminosity of the GRB decreases $\propto \theta^{-2}$, while the degree of linear polarization increases as $P \sim 10^{-2} (\theta/30 ~\mbox{degree})^4$."
 The response of an interferometer (interferometer 1) to the gravitational radiation is eiven bv a linear combination of (wo polarization components my=F4h.+Fyyh. where the antenna patterns Foy and Fy depend on the orientation of the interferometer with respect lo the GW source (e.g. Finn Chernoff 1993).," The response of an interferometer (interferometer $1$ ) to the gravitational radiation is given by a linear combination of two polarization components $m_1=F_{+,1}h_++F_{\times,1} h_\times$ where the antenna patterns $F_{+,1}$ and $F_{\times,1}$ depend on the orientation of the interferometer with respect to the GW source (e.g. Finn Chernoff 1993)."
 We assume that the interferometer arms are of the same length and that thev meet at right angles., We assume that the interferometer arms are of the same length and that they meet at right angles.
 Define a right-handecd coordinate svstem with one interferometer arm along the -axis and the other along the y-axis., Define a right-handed coordinate system with one interferometer arm along the $x$ -axis and the other along the $y$ -axis.
 For simplicity. we assume that the source is in (he direction of the z-axis.," For simplicity, we assume that the source is in the direction of the $z$ -axis."
 Since we can determine the position of a source in the skv by using observations of (he GRD and alterglow. it is possible to generalize the lollowing discussion to the case of a source with an arbitrary sky position.," Since we can determine the position of a source in the sky by using observations of the GRB and afterglow, it is possible to generalize the following discussion to the case of a source with an arbitrary sky position."
 The angular momentum vector of the source (direction of the GRB jet) may be oriented in an arbitrary direction., The angular momentum vector of the source (direction of the GRB jet) may be oriented in an arbitrary direction.
 We assume that the projection of the angular momentum vector to the .7—y plane makes an anele ¢ with the v-axis., We assume that the projection of the angular momentum vector to the $x-y$ plane makes an angle $\zeta$ with the $x$ -axis.
 Wilh these conventions. (he antenna patterns are given by Foy=cos26 and F4=sin26.," With these conventions, the antenna patterns are given by $F_{+,1}=\cos 2\zeta$ and $F_{\times,1}=\sin 2\zeta$."
" To determine the polarization of GWs. one needs a network consisting of at least. two interferometers which have different (1.6, non-parallel) arm orientations."," To determine the polarization of GWs, one needs a network consisting of at least two interferometers which have different (i.e. non-parallel) arm orientations."
 Consider an identical interferometer (interferometer 2) al the same location as interferometer 1. (, Consider an identical interferometer (interferometer 2) at the same location as interferometer 1. (
If we set the arrival time of the GRB signal as the origin of time at each interferometer. we can correct for the actual physical separation so that the (wo interferometers can be always considered to be at the same location).,"If we set the arrival time of the GRB signal as the origin of time at each interferometer, we can correct for the actual physical separation so that the two interferometers can be always considered to be at the same location)."
 We assume that the interferometer 2 is rotated bv an angle 7/4 around (he z-axis with respect to (he interferometer1., We assume that the interferometer 2 is rotated by an angle $-\pi/4$ around the $z$ -axis with respect to the interferometer1.
" The response of the interferometer 2 ds omes=Fosh+Fyohy. and the antenna palters are given by Fo»=—sin26. and ΕνοΞξ60826, "," The response of the interferometer 2 is $m_2=F_{+,2}h_++F_{\times,2} h_\times$, and the antenna patters are given by $F_{+,2}=-\sin 2\zeta,$ and $F_{\times,2}=\cos 2\zeta$."
The detection of the polarization of GWs requires observations wilh a high ratio (SNR) p., The detection of the polarization of GWs requires observations with a high signal-to-noise ratio (SNR) $\rho$.
 A detection is likelier in an optimal case where the wave forms of the poluized components. / and fy. are known.," A detection is likelier in an optimal case where the wave forms of the polarized components, $f_{+}$ and $f_{\times}$, are known."
 We celine the noise-weighted inner product as where ο) denotes the Fourier translorm of the outputs of the two interferometers slt)=i)+n() G— 1.2). nC) is the noise of the interferometers. $5(/f) is the one-siclec," We define the noise-weighted inner product as where $\tilde{s_i}(f)$ denotes the Fourier transform of the outputs of the two interferometers $s_i(t)=m_i(t)+n_i(t)$ $(i=1,2)$ , $n_i(t)$ is the noise of the interferometers, $S_h(f)$ is the one-sided"
three terms were found to be unnecessary descriptors for the data from all of the maps bar the heavily contaminated A band map.,three terms were found to be unnecessary descriptors for the data from all of the maps bar the heavily contaminated $K$ band map.
 This is reassuring as it tells us that the angular power spectrum supplics reliable information about the clata sets., This is reassuring as it tells us that the angular power spectrum supplies reliable information about the data sets.
 We plot in Figures 15 and 16. the results foraddition of the a nonlinear term.," We plot in Figures \ref{fig:reg1} and \ref{fig:reg2}
 the results foraddition of the $z_{j}^2$ non–linear term."
 Curiously. the nonlinearity seen in the A band map is seen at the largest and. very smallest angular separations.," Curiously, the non–linearity seen in the $K$ band map is seen at the largest and very smallest angular separations."
 As has already been noted. this trend is seen in other statistics that we employ.," As has already been noted, this trend is seen in other statistics that we employ."
 The results of applying the skewness. kurtosis and D'Agostino's statistics are shown in Figures 17.0 15. and 19.. respectively.," The results of applying the skewness, kurtosis and D'Agostino's statistics are shown in Figures \ref{fig:skewSPH}, \ref{fig:kurtSPH} and \ref{fig:DSPH}, respectively."
 The results from the ILC map appear consistent with normality., The results from the ILC map appear consistent with normality.
 However. all three statistic behave unusually for the POLL map on scales smaller than (~ 400..," However, all three statistic behave unusually for the TOH map on scales smaller than $\ell\sim 400$ ,."
 Vhis is particularly true of the kurtosis coelIicient where the nonnormality is most evident., This is particularly true of the kurtosis coefficient where the non–normality is most evident.
 It would. be interesting to relate this nonnormality to that already seen in real space., It would be interesting to relate this non–normality to that already seen in real space.
 This could hopefully give us a better handle on the source of the nonCaussianity (whether Galactic or cosmological)., This could hopefully give us a better handle on the source of the non–Gaussianity (whether Galactic or cosmological).
 The univariate transformation results are shown for completeness in Figure 20.., The univariate transformation results are shown for completeness in Figure \ref{fig:uni_transSPH}.
 Llowever. the method appears unreliable as the Gaussian. MC map results are inconsistent with being drawn from à X1 distributions.," However, the method appears unreliable as the Gaussian MC map results are inconsistent with being drawn from a $\chi^2_1$ distributions."
 We feel this unreliability is due to the small values of 9 that make the minimized function shape more complex., We feel this unreliability is due to the small values of $n$ that make the minimized function shape more complex.
 The application of our bivariate skewness and kurtosis statistics are shown in Figures 21. and 22.., The application of our bivariate skewness and kurtosis statistics are shown in Figures \ref{fig:multi_skewSPH} and \ref{fig:multi_kurtosisSPH}. .
 Once again. the ILC map behaviour corresponds to that. of the Gaussian," Once again, the ILC map behaviour corresponds to that of the Gaussian"
Sensitivity may be increased by utilizing a single mask desieu. which increases the trausiussion to5054.,"Sensitivity may be increased by utilizing a single mask design, which increases the transmission to."
 A rotating modulator (RM) (Durouchouxetal.1983:Doadurkeviecius&Ralvs1985) is one such iustruineut developed to image hard x-ray aud gamma-ray photons (teus of keV to MeV).," A rotating modulator (RM) \citep{Durouchoux1983, Dadurkevicius1985} is one such instrument developed to image hard x-ray and gamma-ray photons (tens of keV to MeV)."
 As we have shown previously (Buddenetal.2010b).. the RA Way have some sienificaut sensitivitv advantages over the couunonly-used coded aperture. particularly at high cnereics.," As we have shown previously \citep{BuddenTNS2010}, the RM may have some significant sensitivity advantages over the commonly-used coded aperture, particularly at high energies."
 The RAL cousists of a mask of co-planar parallel slats vpically spaced equidistauce apart rotating above au array of circular detectors (Fig. 1))., The RM consists of a mask of co-planar parallel slats typically spaced equidistance apart rotating above an array of circular detectors (Fig. \ref{fig:rm}) ).
 The transmission of photons from the object scene. $. is modulated iu nue. and so a listory of counts is recorded by each detector.," The transmission of photons from the object scene, $S$, is modulated in time, and so a history of counts is recorded by each detector."
 For a stationary instrument. the recorded data are folded modulo the mask period to produce a count xofile for detector d. Ομ). which iav be described bv if noise is ignored.," For a stationary instrument, the recorded data are folded modulo the mask period to produce a count profile for detector $d$, $O_d(t)$, which may be described by if noise is ignored."
" Pf.i) is the imstrunienut response ""nuetiou. which in this application is a collection of characteristic count rate profiles for poiut sources at all possible scene locations η."," $P_d(t,n)$ is the instrument response function, which in this application is a collection of characteristic count rate profiles for point sources at all possible scene locations $n$."
 huage recoustruction is he technique by which the inverse problem of Eq., Image reconstruction is the technique by which the inverse problem of Eq.
 1 is solved for the object scene ο., \ref{eq:oMatrix} is solved for the object scene $S$.
" Obviously. the modulation xitterus. £P, must be pre-determined aud well-defined."," Obviously, the modulation patterns $P_d$ must be pre-determined and well-defined."
 The ideal techuique for determining the expectel nodulation patterus should be computationally fast. allow for wnconstrained instrament ecometry. account or projection effects and non-uniform attenuation. aud describe the euuulative shadowiug by multiple slats siunultauneouslv.," The ideal technique for determining the expected modulation patterns should be computationally fast, allow for unconstrained instrument geometry, account for projection effects and non-uniform attenuation, and describe the cumulative shadowing by multiple slats simultaneously."
 Brute force Monte Carlo simulations are able to accomplish these tasks. but the computation js fune-consunius.," Brute force Monte Carlo simulations are able to accomplish these tasks, but the computation is time-consuming."
" Durouchouxoetal.(1083) and Doadurkevicius& Balvs(1985) have presented a standi characteristic profile that can ο caleulated analytically, as described iu Sec. 2.."," \cite{Durouchoux1983} and \cite{Dadurkevicius1985} have presented a standard characteristic profile that can be calculated analytically, as described in Sec. \ref{sec:classical}."
 While suitable iu nianv scenarios. this formula imposes tieht constraints on mstriunent eeconietry and is too simplistic to account for non-uniform attenmation or shadow lenethenine.," While suitable in many scenarios, this formula imposes tight constraints on instrument geometry and is too simplistic to account for non-uniform attenuation or shadow lengthening."
 Iu this paper. the previous standard formula is extended aud made iore general. with particular care taken to account for incomplete mask absorption which becomes aportant at hard x-ray and eamuna-rav energies," In this paper, the previous standard formula is extended and made more general, with particular care taken to account for incomplete mask absorption which becomes important at hard x-ray and gamma-ray energies."
 This analysis is part of a program to design a wide-field high resolutiou telescope suitable for hard sx-rav/eanunatray lnaging from a loug-duration balloon or satellite platform (Caimdlayetal.2001:AleCounelletal.2001:Budden 20104).," This analysis is part of a program to design a wide-field high resolution telescope suitable for hard x-ray/gamma-ray imaging from a long-duration balloon or satellite platform \citep{Grindlay2001, McConnell2004, BuddenSORMA2010}."
. To achieve a wide field of view (FOV) aud scusitivity to high euergies Gvhich requires thick mask slats}. a more robust analytical profile is necessary to accurately describe the instrument respouse.," To achieve a wide field of view (FOV) and sensitivity to high energies (which requires thick mask slats), a more robust analytical profile is necessary to accurately describe the instrument response."
" Iu Sec. ον,"," In Sec. \ref{sec:advanced},"
 we present an advanced characteristic profile for the RM that is capable of describing this conrplex imodulatiou pattern aud can be calculated analytically iu a relatively short time., we present an advanced characteristic profile for the RM that is capable of describing this complex modulation pattern and can be calculated analytically in a relatively short time.
 In Sec. L.," In Sec. \ref{sec:results},"
 we show cxamples of count rate profiles generated with the standard aud advanced formulae and with Aloute Carlo sunuulatious. as well as recoustructed nuages.," we show examples of count rate profiles generated with the standard and advanced formulae and with Monte Carlo simulations, as well as reconstructed images."
 The nuage recoustruction techuique has been described elsewhere (Buddenetal.201050)... and has been shown to be capable of providing coded-aperture quality resolution with better detector efficiency at high energies. resolving multiple closely-separated sources and operating in the presence of background.," The image reconstruction technique has been described elsewhere \citep{BuddenTNS2010}, and has been shown to be capable of providing coded-aperture quality resolution with better detector efficiency at high energies, resolving multiple closely-separated sources and operating in the presence of background."
 The standard characteristic formumla for a singleauask RAL was first presented by Durouchoux and examined in ereater detail bv Dacurkevicius&Ralvs (1985)., The standard characteristic formula for a single-mask RM was first presented by \cite{Durouchoux1983} and examined in greater detail by \cite{Dadurkevicius1985}.
. An RM has slat width e. slat spacing b. aud detector ciameter c (Fig. 2)).," An RM has slat width $a$, slat spacing $b$, and detector diameter $c$ (Fig. \ref{fig:geometry}) ),"
 and the mask is suspended a distance £ from the detection lane., and the mask is suspended a distance $L$ from the detection plane.
 The standard formula imposes the coustraüint ο=b c. maxiuizing instrument sensitivitv aud count rate profile contrast.," The standard formula imposes the constraint $a = b = c$ , maximizing instrument sensitivity and count rate profile contrast."
 The assuuiption is made that slats have infinitesimal thickness. but attenuate of the incident photons.," The assumption is made that slats have infinitesimal thickness, but attenuate of the incident photons."
 An attenuation cocfficieu lay be applied to correct for transimission throueh the slats. but clippiug effects. non-uniform attenuation. aud shadow lengthening are ignored.," An attenuation coefficient may be applied to correct for transmission through the slats, but clipping effects, non-uniform attenuation, and shadow lengthening are ignored."
 For the description of both the standard formmla below and the advanced, For the description of both the standard formula below and the advanced
of the main event.,of the main event.
 Finally the distribution of flare energies confirms that the higher the flare energy. the larger the number of subsequent events with high energy.," Finally the distribution of flare energies confirms that the higher the flare energy, the larger the number of subsequent events with high energy."
 The existence of time-energy correlations suggests the possibility of scaling laws relating time with the energy released in a flare., The existence of time-energy correlations suggests the possibility of scaling laws relating time with the energy released in a flare.
 This is a still open question. which could provide interesting insights in the energy storage and release mechanisms at the origin of solar flare occurrence.," This is a still open question, which could provide interesting insights in the energy storage and release mechanisms at the origin of solar flare occurrence."
by MGS835.,by MG85.
 The corresponding dispersion σ of the amplitude shifts Alog.S is σ=0.23. which is slightly smaller than the dispersion of the frequency shifts Alogy (o0= 0.27) and the time shifts Alogf (0= 0.29).," The corresponding dispersion $\sigma$ of the amplitude shifts $\Delta\log{S}$ is $\sigma=0.23$, which is slightly smaller than the dispersion of the frequency shifts $\Delta\log{\nu}$ $\sigma=0.27$ ) and the time shifts $\Delta\log{t}$ $\sigma=0.29$ )."
 The amplitude shiftsAlogS are obviously not correlated with either Aloe» or Alogf (Fig., The amplitude shifts$\Delta\log{S}$ are obviously not correlated with either $\Delta\log{\nu}$ or $\Delta\log{t}$ (Fig.
 daa and b)., \ref{evo1}a a and b).
 This is confirmed by a Spearman rank-order test (Bevington 1969)). which yields that. the observed correlations could occur by chance with a probability of more than On the contrary. the shifts Alosv and Alogf align well along the xf.! line (Fig.," This is confirmed by a Spearman rank-order test (Bevington \cite{B69}) ), which yields that the observed correlations could occur by chance with a probability of more than On the contrary, the shifts $\Delta\log{\nu}$ and $\Delta\log{t}$ align well along the $\nu\!\propto\!t^{-1}$ line (Fig."
 4ec) and the Spearman’s test probability of « In the light-curve approach deseribed above. we model analytically the light. curve of an outburst at. different frequencies and show that the resulting typical flare. is qualitatively in agreement with what is expected by shock models im relativistic jets.," \ref{evo1}c c) and the Spearman's test probability of $<$ In the light-curve approach described above, we model analytically the light curve of an outburst at different frequencies and show that the resulting typical flare is qualitatively in agreement with what is expected by shock models in relativistic jets."
 It is thus of interest to derive from the data the parameters that are relevant to those models., It is thus of interest to derive from the data the parameters that are relevant to those models.
 The shock model of MGS$85 and its generalization. by Valtaoja et al. (1992)), The shock model of MG85 and its generalization by Valtaoja et al. \cite{VTU92}) )
 describe the evolution of the shock by three distinct stages: I) arising phase. 2) a peaking phase and 3) a decliningphase!.," describe the evolution of the shock by three distinct stages: 1) a rising phase, 2) a peaking phase and 3) a declining."
. The three-stage approach presented below is similar to that of Valtaoja et al. (1992)).," The three-stage approach presented below is similar to that of Valtaoja et al. \cite{VTU92}) ),"
 in the sense that its aim is simply to qualitatively describe the observations., in the sense that its aim is simply to qualitatively describe the observations.
 It contains however more parameters in order to include those which are relevant to test the physical model of MG85., It contains however more parameters in order to include those which are relevant to test the physical model of MG85.
 The remarks of Sect., The remarks of Sect.
 3.1 concerning the number of outbursts and the quoted values of the reduced 4? apply equally here., \ref{number} concerning the number of outbursts and the quoted values of the reduced $\chi^2$ apply equally here.
 The self-absorbed synchrotron spectrum emitted by electrons with a powerlaw energy distribution of the form N(E)XE7 can be expressed — by generalizing the homogeneous case (e.g. Pacholezyk 1970:: Stevens et al. 1995)) —, The self-absorbed synchrotron spectrum emitted by electrons with a powerlaw energy distribution of the form $N(E)\!\propto\!E^{-s}$ can be expressed – by generalizing the homogeneous case (e.g. Pacholczyk \cite{P70}; Stevens et al. \cite{SLR95}) ) –
" as where (1/74ji».ΗΕ is equal to the optical depth τ,, at frequency r.", as where $(\nu/\nu_1)^{\alpha\dmrm{thin}-\alpha\dmrm{thick}}$ is equal to the optical depth $\tau_{\nu}$ at frequency $\nu$.
" 5, and » are respectively the flux density and the frequency corresponding to an optical depth of 7,— 1.", $S_1$ and $\nu_1$ are respectively the flux density and the frequency corresponding to an optical depth of $\tau_{\nu}\!=\!1$ .
" At high frequency (772» i4) the medium is optically thin (7,« 1) and the spectrum follows a power law of index oiu,=(s 1)/2."," At high frequency $\nu\!\gg\!\nu_1$ ) the medium is optically thin $\tau_{\nu}\ll 1$ ) and the spectrum follows a power law of index $\alpha\dmrm{thin}\!=\!-(s-1)/2$ ,"
"covers the combination of parameters given in Table 2,, for ages ranging from a few tens to a few thousands of years.","covers the combination of parameters given in Table \ref{table2}, for ages ranging from a few tens to a few thousands of years."
" For each of the above models, we obtained an excitation diagram that can be compared with observations."," For each of the above models, we obtained an excitation diagram that can be compared with observations."
 We describe here the results of the comparisons between observations and our models., We describe here the results of the comparisons between observations and our models.
 A few preliminary remarks about our results are applicable to all the knots and regions analysed., A few preliminary remarks about our results are applicable to all the knots and regions analysed.
" First, the filling factor was considered to be equal to 1 for all transitions; this yields Hz knot sizes consistent with the thickness of the Hz emitting layer of our models."," First, the filling factor was considered to be equal to 1 for all transitions; this yields $_2$ knot sizes consistent with the thickness of the $_2$ emitting layer of our models."
" Secondly, under this assumption, we find that, in all cases, only non-stationary shock models reproduce satisfactorily both the pure rotational and rovibrational (when available) parts of the Ho excitation diagram, a result that has already been established in similar studies of other objects (e.g. HH54, ?,, or L1157-B1, see G08b)."," Secondly, under this assumption, we find that, in all cases, only non-stationary shock models reproduce satisfactorily both the pure rotational and rovibrational (when available) parts of the $_2$ excitation diagram, a result that has already been established in similar studies of other objects (e.g. HH54, \citealt{Giannini06}, or L1157-B1, see G08b)."
" Finally, the departure of the observational points from a straight line in the purely rotational parts of the excitation diagram (“saw-tooth’ pattern, which is weak in our cases) is indicative of deviations of the ortho-to-para ratio from the high-temperature LTE value of 3.0."," Finally, the departure of the observational points from a straight line in the purely rotational parts of the excitation diagram saw-tooth' pattern, which is weak in our cases) is indicative of deviations of the ortho-to-para ratio from the high-temperature LTE value of 3.0."
 Such a pattern can be reproduced in the models by slightly modifying the value of the initial ortho-to-para ratio; but fine-tuning this parameter is not a major consideration in the context of the present study., Such a pattern can be reproduced in the models by slightly modifying the value of the initial ortho-to-para ratio; but fine-tuning this parameter is not a major consideration in the context of the present study.
" We give prior focus to the Herbig-Haro objects, for which we can benefit from H» excitation diagrams combining pure rotational and rovibrational transitions, and to the so-called “SiO knot’, for which future observations of SiO rotational transitions are expected to provide further constraints on our shock models."," We give prior focus to the Herbig-Haro objects, for which we can benefit from $_2$ excitation diagrams combining pure rotational and rovibrational transitions, and to the so-called SiO knot', for which future observations of SiO rotational transitions are expected to provide further constraints on our shock models."
 Tables Al and A2 present a summary of the restrictions that comparisons between observational and computed excitation diagrams place on the shock model parameters., Tables \ref{tablea1} and \ref{tablea2} present a summary of the restrictions that comparisons between observational and computed excitation diagrams place on the shock model parameters.
" In these tables, the values of the pre-shock density, magnetic-field parameter and shock velocity are indicated for our best models, as well as the age range for which the excitation diagrams are in satisfactory lagreement with the observations."," In these tables, the values of the pre-shock density, magnetic-field parameter and shock velocity are indicated for our best models, as well as the age range for which the excitation diagrams are in satisfactory lagreement with the observations."
" In the case of the Herbig-Haro objects B and HH321A/B,the best-fitting models are non-stationary, and the pre-shock density is constrained to 10* cm?, which is in agreement with the average density of the whole globule (9x10? cm?), as determined by ?.."," In the case of the Herbig-Haro objects B and B,the best-fitting models are non-stationary, and the pre-shock density is constrained to $^4$ $^{-3}$, which is in agreement with the average density of the whole globule $\times$ $^{3}$ $^{-3}$ ), as determined by \citet{Bourke97}. ."
" Shock velocities in the range 15-35 km s! yield good fits to the observations,"," Shock velocities in the range 15–35 km $^{-1}$ yield good fits to the observations,"
Combined Array for Research in Millimeter-wave Astronomy and the Expanded Very Large Array(EVLA)**.,Combined Array for Research in Millimeter-wave Astronomy and the Expanded Very Large Array.
. A few days later low frequency observations were undertaken at Westerbork Svuthesis Radio Telescope(WSRT)74., A few days later low frequency observations were undertaken at Westerbork Synthesis Radio Telescope.
. X-ray observations were obtained with theSwift and Chandra observatories., X-ray observations were obtained with the and Chandra observatories.
 The early optical lisht curve of ΕΕΚ shows an extra-ordinary good fit to that. expected [rom an explocing star (flux proportional to exponential of square of time)., The early optical light curve of PTF11kly shows an extra-ordinary good fit to that expected from an exploding star (flux proportional to exponential of square of time).
 As a result the birth ol the supernova can be accurately timed to a fraction an hour: UT 2011 August 23.69 (Nugent et al. 20113)., As a result the birth of the supernova can be accurately timed to a fraction an hour: UT 2011 August 23.69 (Nugent et al. \nocite{Nugent11}) ).
 Our first observations at both radio and. X-ray. bands were taken just over a day. after (he explosion., Our first observations at both radio and X-ray bands were taken just over a day after the explosion.
 The log of observations and the associated details can be found in Table L.., The log of observations and the associated details can be found in Table \ref{tab:RadioLog}.
 Our CARALA and EVLA observations include the earliest search for radio emission in em-wave and mm-wave bands and subsequent observations include a very sensitive search in the 21-cm band obtained at the WSRT (see Figure 1))., Our CARMA and EVLA observations include the earliest search for radio emission in cm-wave and mm-wave bands and subsequent observations include a very sensitive search in the 21-cm band obtained at the WSRT (see Figure \ref{fig:WSRTImage}) ).
 Following our first EVLA and CARAIA observations. additional data was," Following our first EVLA and CARMA observations, additional data was"
,.
2002).. Agolοἱal.(2002) have suggested that current. microlensing data indicate the presence of a significant population of intermediate mass black holes roaming the Galactic disk: interferometers will be able to confirm or reject this possibility., \citet{agol} have suggested that current microlensing data indicate the presence of a significant population of intermediate mass black holes roaming the Galactic disk; interferometers will be able to confirm or reject this possibility.
 In this paper. we have focused on ground-based interferometers. however our results apply also to space-basecl interferometers like the Space Interferometry Mission. (SIM).," In this paper, we have focused on ground-based interferometers, however our results apply also to space-based interferometers like the Space Interferometry Mission (SIM)."
 SIM is primarily an astrometric insirument. however it can also measure fringe visibilities.," SIM is primarily an astrometric instrument, however it can also measure fringe visibilities."
 Nominally. the target precision expected for SIM is in V? (M. Shao 2002. ccomm.).," Nominally, the target precision expected for SIM is in $V^2$ (M. Shao 2002, comm.)."
 SIMs baseline is 10m. and typical wavelengths are A2Q.6jan. giving resolution of about 12 mas.," SIM's baseline is 10m, and typical wavelengths are $\lambda\approx0.6\mu$ m, giving resolution of about 12 mas."
 Hence. SIM can determine 65 with 10 from visibility alone. entirely independently ol the astrometric determination. for events with 8j>0.414 mas.," Hence, SIM can determine $\thetaE$ with $>10$ from visibility alone, entirely independently of the astrometric determination, for events with $\thetaE>0.44$ mas."
 From Figure 5 we see that (his comprises a large Iraction of the events., From Figure \ref{events} we see that this comprises a large fraction of the events.
 The visibility measurements come for free with the astrometric measurements. ancl should significantly increase the precision of SIAL mass nmeasurenienis. as long as effects such as crowding do not pose too great an obstacle.," The visibility measurements come for free with the astrometric measurements, and should significantly increase the precision of SIM mass measurements, as long as effects such as crowding do not pose too great an obstacle."
" In addition to measuring je. SIM also determines πιο, the lens parallax. bv measuring the time of the peak of the photometric lighteurve."," In addition to measuring $\thetaE$, SIM also determines $\piE$, the lens parallax, by measuring the time of the peak of the photometric lightcurve."
 Since the peak of the visibilitv signal coincides with the peak of the photometric signal. SIM visibility measurements could also be used to determine zi.," Since the peak of the visibility signal coincides with the peak of the photometric signal, SIM visibility measurements could also be used to determine $\piE$."
 However. since the variation in 1—V? is so shallow near the peak this method nav not prove to be as precise as ordinarv photometric parallax measurement.," However, since the variation in $1-V^2$ is so shallow near the peak this method may not prove to be as precise as ordinary photometric parallax measurement."
 One of the most (potentially) exciting prospects is a topic we have not discussed in (his paper. binary microlensing.," One of the most (potentially) exciting prospects is a topic we have not discussed in this paper, binary microlensing."
 For a binary lens svstem. complicated caustic structures can arise. leading in favorable cases to the production of 5 images of the source.," For a binary lens system, complicated caustic structures can arise, leading in favorable cases to the production of 5 images of the source."
 For a spectacular example of this. see Anetal.(2002).," For a spectacular example of this, see \citet{an}."
. During caustic erossings. the magnification can get exceptionally large. e.g. factors of 30. making these events bright enough to observe with interlerometers.," During caustic crossings, the magnification can get exceptionally large, e.g. factors of 30, making these events bright enough to observe with interferometers."
 The five images are currently unresolvable from each other. however VLTI and ουκ offer the prospect of imaging the mulü-image pattern.," The five images are currently unresolvable from each other, however VLTI and Keck offer the prospect of imaging the multi-image pattern."
 With a multi-aperture svslem (required [or closure phase). one obtains several visibility. measurements aud one closure phase at the same time. possibly allowing the reconstruction of complex events such as causlic crossings.," With a multi-aperture system (required for closure phase), one obtains several visibility measurements and one closure phase at the same time, possibly allowing the reconstruction of complex events such as caustic crossings."
the passage of a shock. the actual collapse time will be somewhat longer than the free-fall lime due (ο turbulence within the cloud.,"the passage of a shock, the actual collapse time will be somewhat longer than the free-fall time due to turbulence within the cloud."
 In this work. we use the free-fall time as a useful under-estimate of the (vpical collapse time for a molecular cloud undergoing quiescent star formation.," In this work, we use the free-fall time as a useful under-estimate of the typical collapse time for a molecular cloud undergoing quiescent star formation."
 We can relate the Iree-Iall time of a collapsing core to the number density of hydrogen nuclei as follows: where nyΠΠ)2n(15)) is the local number clensity of hydrogen nuclei in oE (?).., We can relate the free-fall time of a collapsing core to the number density of hydrogen nuclei as follows: where $n_H=(n(H)+2n(H_{2}))$ is the local number density of hydrogen nuclei in $^{-3}$ \citep{DysonWilliams:ISM}.
 In order for a cloud of eas to continue to collapse and form a star. (e cloud must continue to lose energy. otherwise. the high temperatures and densities within the cloud would halt collapse before (he formation of a star could occur.," In order for a cloud of gas to continue to collapse and form a star, the cloud must continue to lose energy, otherwise, the high temperatures and densities within the cloud would halt collapse before the formation of a star could occur."
 Energy loss occurs via different. atomic and molecular transitions. while at high densities. collisions between hot gas and cooler dust erains can also be ellicient in cooling the gas.," Energy loss occurs via different atomic and molecular transitions, while at high densities, collisions between hot gas and cooler dust grains can also be efficient in cooling the gas."
 The rate al which energy is lost depends on the cooling [unction. .V;/ erg em. oE I. which in turn depends on the abundance of various different molecular coolants.," The rate at which energy is lost depends on the cooling function, $\Lambda_{tot}/$ erg cm $^{-3}$ $^{-1}$, which in turn depends on the abundance of various different molecular coolants."
 The total cooling time is (hen given bv: where fis the Doltzmanun constant. Z'/is the gas temperature and £ is the metallicity in units of solar metallicity.," The total cooling time is then given by: where $k$ is the Boltzmann constant, $T$ is the gas temperature and $\xi$ is the metallicity in units of solar metallicity."
" The quantitv. 27,42 represents the total energy. both thermal and internal. of the svstem and £4 is the total interstellar cooling rate."," The quantity, $\frac{5}{2}n_HkT$ represents the total energy, both thermal and internal, of the system and $\xi\Lambda_{tot}$ is the total interstellar cooling rate."
 Note that our chemical models described in § 3.. do not include the effects of removing molecular coolants [rom the gas pliase by [reeze-out.," Note that our chemical models described in $\S$ \ref{sec:model}, do not include the effects of removing molecular coolants from the gas phase by freeze-out."
 If [reeze-out occurs. (he cooling timescale will be underestimated in regions of parameter space where (here is significant ice-formation.," If freeze-out occurs, the cooling timescale will be underestimated in regions of parameter space where there is significant ice-formation."
 A detailed caleulation of the rate per unit volume at which a species freezes out can be, A detailed calculation of the rate per unit volume at which a species freezes out can be
οποιον density is assuined to be coustaut. p=const. (18) is reduced to the following equation: thus. «is not necessarily coustaut.,"energy density is assumed to be constant, $\rho=const.$, (18) is reduced to the following equation: thus, $w$ is not necessarily constant."
 Similarly. if l. (18) is reduced to the following equation: thus. the euergv density is not necessarily coustaut cither.," Similarly, if $w=-1$ , (18) is reduced to the following equation: thus, the energy density is not necessarily constant either."
 Using the energsvauonientuni tensor paranetrzed iu (17). the auisotropy enerev density can be written in ternis of the directional EoS parameters as follows: Note that the behavior of the anisotropy energy deusity is not oulv determined by the Aj; coustauts aud the volune of the universe as in (16). but also bv the directional EoS poarunueters (a. wy aud wi) aud the οποιον density of the fluid (p) via the integral term.," Using the energy-momentum tensor parametrized in (17), the anisotropy energy density can be written in terms of the directional EoS parameters as follows: Note that the behavior of the anisotropy energy density is not only determined by the $\lambda_{j}$ constants and the volume of the universe as in (16), but also by the directional EoS parameters $w_{x}$, $w_{y}$ and $w_{z}$ ) and the energy density of the fluid $\rho$ ) via the integral term."
 We imav write down the ecneralized Eriediuaun equation bv considering (13) aud (11) and define the effective energy density pig. which determines the voluuetiic expansion rate of the uulverse. From the definition of the mean IIubble parameter eiven in (9) the volume of the uuiverse can be obtained as follows: where eq is a positive constant of integration.," We may write down the generalized Friedmann equation by considering (13) and (14) and define the effective energy density $\rho_{\textnormal{ef}}$, which determines the volumetric expansion rate of the universe, From the definition of the mean Hubble parameter given in (9) the volume of the universe can be obtained as follows: where $c_{1}$ is a positive constant of integration."
 Thus. using (22) in (23) the volume of the universe cau be obtained m terms of the effective energy. density. According to this. the conditiou for the de Sitter volumetric expansion does uot correspond to a constaut energv density of the finid but to a constant effective euergv density. Le. pog=p|ps const. which leads to Obviously. the universe exhibits de Sitter expansion when p; is null iu the presence of a positive cosinological constant. c.. in the preseuce of conventional vacuna energv that wiclds a constant enerev deusity.," Thus, using (22) in (23) the volume of the universe can be obtained in terms of the effective energy density, According to this, the condition for the de Sitter volumetric expansion does not correspond to a constant energy density of the fluid but to a constant effective energy density, i.e., $\rho_{\textnormal{ef}}=\rho+\rho_{\beta}=const.$ , which leads to Obviously, the universe exhibits de Sitter expansion when $\rho_{\beta}$ is null in the presence of a positive cosmological constant, i.e., in the presence of conventional vacuum energy that wields a constant energy density."
 However. if ps is uot null then pg is uot constaut. but a decreasing function of the cosmüc £ as long as the mniverse expands. since p;xV.2.," However, if $\rho_{\beta}$ is not null, then $\rho_{\textnormal{ef}}$ is not constant, but a decreasing function of the cosmic $t$ as long as the universe expands, since $\rho_{\beta}\propto V^{-2}$."
 Thus; as also shown w DBoeesluuu(1991).. the behavior of Biauchi type-l inflationary solutions cannot be of the pure de Sitter ype. even when p=p. but are of power-law type.," Thus, as also shown by \cite{beesham94}, , the behavior of Bianchi type-I inflationary solutions cannot be of the pure de Sitter type, even when $p=-\rho$, but are of power-law type."
 What if we consider the anisotropy energv density ogether with a mixture of the conventional vacumu energv and perfect fluids?, What if we consider the anisotropy energy density together with a mixture of the conventional vacuum energy and perfect fluids?
 The conventional perfect Huids. Lo. radiation. pressureless matter etc.," The conventional perfect fluids, i.e., radiation, pressureless matter etc."
 can be deseribed by an EoS paramcter in the form of p=wp and their euergy deusities change as V.LUCI, can be described by an EoS parameter in the form of $p=w\rho$ and their energy densities change as $V^{-(1+w)}$.
 Thias. the cosmological constant will be dominant as V.»|x. provided that (e2/—1. and according to the cosmic hair theorem for Eiustcin gravitv introduced by Wald(1983).. the universe will expoucutially evolve toward the de Sitter universe.," Thus, the cosmological constant will be dominant as $V\rightarrow+\infty$, provided that $w>-1$, and according to the cosmic no-hair theorem for Einstein gravity introduced by \cite{wald83}, the universe will exponentially evolve toward the de Sitter universe."
 Iu other words. Bianchi tvpe-I models iu the preseuce of a positive cosinological constant isotropize and their volumetric expuiuisiou rates approach de Sitter expausionas V.»|x.," In other words, Bianchi type-I models in the presence of a positive cosmological constant isotropize and their volumetric expansion rates approach de Sitter expansion as $V\rightarrow+\infty$."
 On the other hand. according to (16). no matter how small the anisotropy enerev deusitv is. compared with the other sources. pa will eventually dominate amy perfect Suid and govern the dynamics of the expansion in the very carly evolution of the universe as >0. provided that coc Ἐν nes the universe will approximate the hNasner vacuum solution.," On the other hand, according to (16), no matter how small the anisotropy energy density is, compared with the other sources, $\rho_{\beta}$ will eventually dominate any perfect fluid and govern the dynamics of the expansion in the very early evolution of the universe as $V\rightarrow 0$, provided that $w<1$ , i.e., the universe will approximate the Kasner vacuum solution."
 However. these cases are not implied ouce the implicitly assiued isotropy of the vacua euerev is relaxed.," However, these cases are not implied once the implicitly assumed isotropy of the vacuum energy is relaxed."
 This is because. according to (21). p; does uot have to be proportional with V.7? in case of an auisotropie fiuid. aud it may not diverse as V.$0 and/or V»| x.Qur final remark in this section coucerus the conditionlfor isotropization of a Biauchi type-I universe that exhibits de Sitter volumetric expansion.," This is because, according to (21), $\rho_{\beta}$ does not have to be proportional with $V^{-2}$ in case of an anisotropic fluid, and it may not diverge as $V\rightarrow 0$ and/or $V\rightarrow+\infty$ .Our final remark in this section concerns the conditionfor isotropization of a Bianchi type-I universe that exhibits de Sitter volumetric expansion."
 The condition for isotropization mentioned in Sect., The condition for isotropization mentioned in Sect.
 2.3 can be rewritteu as follows: ρ> Gas f>|x since JT is coustaut. aud this leads to py< Qas foo»|x. since ps>0 by definition.," 2.3 can be rewritten as follows: $\rho_{\beta}\rightarrow0$ as $t\rightarrow+\infty$ since $H$ is constant, and this leads to $\dot{\rho_{\beta}}<0$ as $t\rightarrow+\infty$, since $\rho_{\beta}>0$ by definition."
 This. in addition. leads to p>ϐ as fooqox from ps|p=0.," This, in addition, leads to $\dot{\rho}>0$ as $t\rightarrow+\infty$ from $\dot{\rho_{\beta}}+\dot{\rho}=0$."
 Thus. the condition for isotropization in the Bianchi tvpe-I models that exhibit de Sitter volumetric expansion imposes ou the fluid the constraint to behave like a phantom cuerev. be. to exhibit an increasing euergv density as V increases.," Thus, the condition for isotropization in the Bianchi type-I models that exhibit de Sitter volumetric expansion imposes on the fluid the constraint to behave like a phantom energy, i.e., to exhibit an increasing energy density as $V$ increases."
 Ou the other hand. if the energv densityof the fiuid is coustaut(p= const). then the anisotropy energy density is alsocoustaut (p;— const).," On the other hand, if the energy densityof the fluid is constant$\rho=const.$ ), then the anisotropy energy density is alsoconstant $\rho_{\beta}=const.$ )."
 If p«0. then po 0: thus for models in which the energy. deusity of the fluid is decreasing. the anisotropy cucrey deusity," If $\dot{\rho}<0$, then $\dot{\rho_{\beta}}>0$ ; thus for models in which the energy density of the fluid is decreasing, the anisotropy energy density"
We calculate Z and § for source bundles with radii Ay = 0.01.0.02.0.05.0.1.0.2.0.5.1 and 2.,"We calculate ${\cal I}$ and ${\cal S}$ for source bundles with radii $\Delta y$ = $0.01, 0.02, 0.05, 0.1, 0.2, 0.5, 1$ and 2."
 In Fig. S..," In Fig. \ref{fig:variation},"
 we plot Z and Sas [functions of the image bundle impact parameter. p.," we plot ${\cal I}$ and ${\cal S}$ as functions of the image bundle impact parameter, $x_c$."
 We select Z2 tand S>4 as indicative that the distortion is significant: this was confirmed by exe using plots similar to Fig. 7.., We select ${\cal I} \geq -4$ and ${\cal S} \geq -4$ as indicative that the distortion is significant; this was confirmed by eye using plots similar to Fig. \ref{fig:faceplot}.
 For higher values of Z and S. there were clear dillerences between sources A and C. and images D and D. We find for comparison of source bundle shapes A and ο 1) and for comparison of image bundle shapes Band D (Z> 1).," For higher values of ${\cal I}$ and ${\cal S}$, there were clear differences between sources A and C, and images B and D. We find for comparison of source bundle shapes A and C ${\cal S} \geq -4$ ) and for comparison of image bundle shapes B and D ${\cal I} \geq -4$ )."
 Next. we consider a more quantitative approach. based on comparing buncdle ellipticities. in à manner comparable to the standard analysis for examining weak lensing-induced shear.," Next, we consider a more quantitative approach, based on comparing bundle ellipticities, in a manner comparable to the standard analysis for examining weak lensing-induced shear."
 Defining quadrupole terms relative to the bundle centroid we consider a complex cllipticity (c.g. Schneider. 2005) ol the form: which has norm: We define error terms: for comparison between ellipticities. and determine the image plane impact parameter. à. at which a source with raclius. Ay. Hirst exceeds do=14. 54 and LOM.," Defining quadrupole terms relative to the bundle centroid we consider a complex ellipticity (e.g. Schneider 2005) of the form: which has norm: We define error terms: for comparison between ellipticities, and determine the image plane impact parameter, $x_c$ , at which a source with radius, $\Delta y$ , first exceeds $E = 1\%$, $5\%$ and $10\%$."
 Since a circular source has [yf=0. we use elliptical sources with axis ratios ία=0.5. 0.9 and 0.99.," Since a circular source has $\vert \chi \vert = 0$, we use elliptical sources with axis ratios $b/a = 0.8$, $0.9$ and $0.99$."
 We consider the two cases where the semi-major axis is aligned tangentially or raclially to the Einstein radius. computing these limits ancl also the average (based. on original data values) of these two orientations.," We consider the two cases where the semi-major axis is aligned tangentially or radially to the Einstein radius, computing these limits and also the average (based on original data values) of these two orientations."
" Plotting results as log)(r,1) versus log),Ay. see Fig. 9.."," Plotting results as $\log_{10} (x_c - 1)$ versus $\log_{10} \Delta y$ , see Fig. \ref{fig:logger},"
 we see a relationship that is highly suggestive of a functional form: We perform a least-squares fit in log-log space to obtain the parameters ο and n., we see a relationship that is highly suggestive of a functional form: We perform a least-squares fit in log-log space to obtain the parameters $\epsilon$ and $n$.
 Results of these fits are presented. in ‘Table 2 in all cases. the caleulatec Pearson cocllicient is rol0.994. indicative that equation (55)) is an appropriate functional form.," Results of these fits are presented in Table \ref{tbl:lsf1} – in all cases, the calculated Pearson coefficient is $r > 0.994$, indicative that equation \ref{eqn:funcform}) ) is an appropriate functional form."
 There is variation in the fitted parameters based on the chosen source axis ratio., There is variation in the fitted parameters based on the chosen source axis ratio.
 This is not suprising. as the tical gravitational field. across the resultant. image depends on the relative separations ancl orientation of individual image ravs [rom the lens (ο," This is not suprising, as the tidal gravitational field across the resultant image depends on the relative separations and orientation of individual image rays from the lens (c.f."
 with discussion on orientation of image bundles in section 3.4)., with discussion on orientation of image bundles in section 3.4).
 As a best estimate. we average over the three chosen axis ratios [or Eae=5%. to obtain the second-to-LIast row in Table 2..," As a best estimate, we average over the three chosen axis ratios for $E_{AC} = 5\%$, to obtain the second-to-last row in Table \ref{tbl:lsf1}."
 We propose the following interpretation: if we see an image that looks like D. and we use the IHlexion formalism to determine what the source would look like. we would be wrong (error of 25% in source ellipticity) unless: We refer to the boundary. defined by this expression as the start of the “Hexion zone” (blue line and both grev- regions in Fig. 10)).," We propose the following interpretation: if we see an image that looks like B, and we use the flexion formalism to determine what the source would look like, we would be wrong (error of $\gtrsim 5\%$ in source ellipticity) unless: We refer to the boundary defined by this expression as the start of the “flexion zone” (blue line and both grey-shaded regions in Fig. \ref{fig:crossover}) )."
 For image impact parameters closer to the lens thanthislimit (white region in Fig. 10)).," For image impact parameters closer to the lens thanthislimit (white region in Fig.\ref{fig:crossover}) ),"
 the seconc-order “Taylor series. approximation given bv equation (6)) is not suflicienthy accurate when applied. to, the second-order Taylor series approximation given by equation \ref{eqn:Dijk}) ) is not sufficiently accurate when applied to
 (28)).,\ref{eq:casa}) ).
 That they are just specific cases of what is at core a very simple and elegant equation is a point perhaps so obvious that some authors do not bother noting it. but it cannot be stressed enough!," That they are just specific cases of what is at core a very simple and elegant equation is a point perhaps so obvious that some authors do not bother noting it, but it cannot be stressed enough!"
 The second pitfall is that an equation like (28)). when implemented in software. can be both too specific. and insufficiently. flexible. (," The second pitfall is that an equation like \ref{eq:casa}) ), when implemented in software, can be both too specific, and insufficiently flexible. ("
"Note that the CASA implementation specifies both the time/frequeney behaviour. and the form of the Jones terms. e.g. G is diagonal and variable in time. B is diagonal and variable in frequency. D has a specific ""leakage"" form. etc.)","Note that the CASA implementation specifies both the time/frequency behaviour, and the form of the Jones terms, e.g. $\jones{G}{}$ is diagonal and variable in time, $\jones{B}{}$ is diagonal and variable in frequency, $\jones{D}{}$ has a specific “leakage” form, etc.)"
 For instance. the calibration described in Paper IIT (?) cannot be done in CASA. despite using an ostensibly much simpler form of the RIME. because it includes a Jones term that was not anticipated in the CASA design.," For instance, the calibration described in Paper III \citep{RRIME3} cannot be done in CASA, despite using an ostensibly much simpler form of the RIME, because it includes a Jones term that was not anticipated in the CASA design."
" A second major virtue of the RIME is its ability to describe different propagation effects: this is immediately compromised if only a specific and limited set of these 1s chosen for implementation,", A second major virtue of the RIME is its ability to describe different propagation effects; this is immediately compromised if only a specific and limited set of these is chosen for implementation.
 A final pitfall of the Jones-specific. view 1s that it tends to stereotype approaches to calibration., A final pitfall of the Jones-specific view is that it tends to stereotype approaches to calibration.
" Equation (28)) is a huge improvement on the approaches of older software systems. but in the end it is just some model of an interferometer that happens to work well enough for ""classically-designed"" instruments such as the VLA and WSRT. in their most common regimes."," Equation \ref{eq:casa}) ) is a huge improvement on the approaches of older software systems, but in the end it is just some model of an interferometer that happens to work well enough for “classically-designed” instruments such as the VLA and WSRT, in their most common regimes."
" It is universally true that polarization effects can be completely described by à direction-independent leakage matrix (D). or bandpass by B, — it just happens to be a practical first-order model. which completely breaks down for a new instrument such as LOFAR. where e.g. ""leakage"" is strongly direction-dependent."," It is universally true that polarization effects can be completely described by a direction-independent leakage matrix $\jones{D}{p}$ ), or bandpass by $\jones{B}{p}$ – it just happens to be a practical first-order model, which completely breaks down for a new instrument such as LOFAR, where e.g. “leakage” is strongly direction-dependent."
 In fact. even WSRT results can be improved by departing from this model. as Paper HI (2) will show.," In fact, even WSRT results can be improved by departing from this model, as Paper III \citep{RRIME3} will show."
 We must therefore take care that our thinking about calibration does not fall into a rut marked out by a specific series of Jones terms., We must therefore take care that our thinking about calibration does not fall into a rut marked out by a specific series of Jones terms.
 In Sect. 1.. ," In Sect. \ref{sec:derivation}, ,"
I mentioned that the RIME holds in any coordinate system., I mentioned that the RIME holds in any coordinate system.
 9? briefly discussed coordinate transforms in. this context. but a few additional words on the subject are required.," \citet{ME1} briefly discussed coordinate transforms in this context, but a few additional words on the subject are required."
 Field vectors e and Jones matrices J may be represented [by a particular set of complex values] in any coordinate system. by picking a pair of complex basis vectors in the plane orthogonal to the direction of propagation.," Field vectors $\vec e$ and Jones matrices $\jones{J}{}$ may be represented [by a particular set of complex values] in any coordinate system, by picking a pair of complex basis vectors in the plane orthogonal to the direction of propagation."
 | have usec an orthonormal xv system until now., I have used an orthonormal $xy$ system until now.
" Another useful system is that of circular polarization coordinates. 77. whose basis vectors (represented in the xv system) are e,=-&(1.—i) anc e;=(l.i."," Another useful system is that of circular polarization coordinates $rl$, whose basis vectors (represented in the $xy$ system) are $\vec e_r=\frac{1}{\sqrt{2}}(1,-i)$ and $\vec e_l=\frac{1}{\sqrt{2}}(1,i)$."
 Any other pair of basis vectors may of course be used., Any other pair of basis vectors may of course be used.
" In general. for any two coordinate systems S and T. there will be a corresponding 2x T. such that ej;=Tes. where es and e, represent the same vector in the S and T coordinate systems."," In general, for any two coordinate systems S and T, there will be a corresponding $2\times2$ $\jones{T}{}$, such that $\vec e_\mathrm{T}=\jones{T}{} \vec e_\mathrm{S}$, where $\vec e_\mathrm{S}$ and $\vec e_\mathrm{T}$ represent the same vector in the S and T coordinate systems."
 Likewise. the representation of the linear operator J transforms as Jy=TJST. while the brightness matrix B. (or indeed any cohereney matrix) transforms as By=ΤΒςΤΗ. Of particular importance is the matrix for conversion from linear to circularly polarized coordinates.," Likewise, the representation of the linear operator $\jones{J}{}$ transforms as $\jones{J}{\mathrm{T}}=\jones{T}{} \jones{J}{\mathrm{S}} \jonesinv{T}{}$, while the brightness matrix $\coh{B}{}$ (or indeed any coherency matrix) transforms as $\coh{B}{\mathrm{T}}=\jones{T}{} \coh{B}{\mathrm{S}} \jonesT{T}{}.$ Of particular importance is the matrix for conversion from linear to circularly polarized coordinates."
" This matrix is commonly designated as H (being the mathematical equivalent of an electronicAivbrid sometimes found in antenna receivers): Consequently. the brightness matrix B. when represented in circular polarization coordinates. has the following form (Ul use the indices ""Oo"" and ""+"" where necessary to disambiguate between circular and linear representations): While EMF vectors and Jones matrices may be represented using an arbitrary. basis. the. receptor voltages we actually measure are specific numbers."," This matrix is commonly designated as $\jones{H}{}$ (being the mathematical equivalent of an electronic sometimes found in antenna receivers): Consequently, the brightness matrix $\coh{B}{}$, when represented in circular polarization coordinates, has the following form (I'll use the indices $\odot$ ” and $+$ ” where necessary to disambiguate between circular and linear representations): While EMF vectors and Jones matrices may be represented using an arbitrary basis, the receptor voltages we actually measure are specific numbers."
 The voltage measurement process thus implies a coordinate system. Le. circular for circular receptors. and linear for linear receptors.," The voltage measurement process thus implies a coordinate system, i.e. circular for circular receptors, and linear for linear receptors."
 It is of course possible to convert measured data into a different coordinate frame after the fact., It is of course possible to convert measured data into a different coordinate frame after the fact.
 It is also perfectly possible. and indeed may be desirable. to mix coordinate systems within the RIME. by inserting appropriate coordinate conversion matrices into the Jones chain.," It is also perfectly possible, and indeed may be desirable, to mix coordinate systems within the RIME, by inserting appropriate coordinate conversion matrices into the Jones chain."
" A commonly encountered assumption is that à “VLA RIME"" must be written down in circular coordinates and à ""WSRT RIME'"" in linear. but this is by no means a fundamental requirement!"," A commonly encountered assumption is that a “VLA RIME” must be written down in circular coordinates and a “WSRT RIME” in linear, but this is by no means a fundamental requirement!"
 We're free to express part of the signal propagation. chain in one coordinate. frame. then insert conversion. matrices. at the appropriate place in the equation to switch to a different coordinate frame.," We're free to express part of the signal propagation chain in one coordinate frame, then insert conversion matrices at the appropriate place in the equation to switch to a different coordinate frame."
 In the onion form of the RIME (Eq. 9)).," In the onion form of the RIME (Eq. \ref{eq:me0-onion}) ),"
 this corresponds to à change of coordinate systems as we go from one layer of the onion to another., this corresponds to a change of coordinate systems as we go from one layer of the onion to another.
 For example: One reason to consider the use of mixed coordinate systems is the opportunity to optimize the representation of particular physical effects., For example: One reason to consider the use of mixed coordinate systems is the opportunity to optimize the representation of particular physical effects.
 As an example. a rotation in the xv frame (e.g. tonospheric Faraday rotation. or parallactic angle) is represented by a diagonal matrix in the 77 frame.," As an example, a rotation in the $xy$ frame (e.g. ionospheric Faraday rotation, or parallactic angle) is represented by a diagonal matrix in the $rl$ frame."
 If the observed field has no intrinsic. linear polarization. the Β.. matrix is also diagonal.," If the observed field has no intrinsic linear polarization, the $\coh{B}{\odot}$ matrix is also diagonal."
 If a part of the RIME ts known to contain diagonal matrices only. their product can be evaluated with significant computational savings (compared to the full 2x matrix regime).," If a part of the RIME is known to contain diagonal matrices only, their product can be evaluated with significant computational savings (compared to the full $2\times2$ matrix regime)."
 On the other hand. if the instrument 1s using linear receptors. then receiver gains (G) should be expressed in the linear frame. lest calibrating them become extremely awkward.," On the other hand, if the instrument is using linear receptors, then receiver gains $\jones{G}{}$ ) should be expressed in the linear frame, lest calibrating them become extremely awkward."
 We should therefore implement the RIME somewhat like the above equation. with the appropriate H matrices inserted as “late” in the chain as possible. so that only the minimum amount of computation is done for the full 2x2 case.," We should therefore implement the RIME somewhat like the above equation, with the appropriate $\jones{H}{}$ matrices inserted as “late” in the chain as possible, so that only the minimum amount of computation is done for the full $2\times2$ case."
 This approach is not yet exploited by any existing software. but perhaps it should be.," This approach is not yet exploited by any existing software, but perhaps it should be."
 In particular. the MeqTrees system (?) automatically optimizes internal calculations when only diagonal matrices are in play. and would provide a suitable vehicle for exploring thistechnique.," In particular, the MeqTrees system \citep{meqtrees} automatically optimizes internal calculations when only diagonal matrices are in play, and would provide a suitable vehicle for exploring thistechnique."
"that have environments of a given percentage rank plotted against the host halo mass, for four different nearest neighbour-based techniques, with the number of neighbours increasing from left to right.","that have environments of a given percentage rank plotted against the host halo mass, for four different nearest neighbour-based techniques, with the number of neighbours increasing from left to right."
" These are: the 3rd nearest neighbour density in three dimensions, the surface density for the projected 7th nearest neighbour, the three dimensional density using a 10 neighbour Bayesian metric, and the smooth kernel three dimensional density using 64 neighbours."," These are: the 3rd nearest neighbour density in three dimensions, the surface density for the projected 7th nearest neighbour, the three dimensional density using a 10 neighbour Bayesian metric, and the smooth kernel three dimensional density using 64 neighbours."
" The most noticeable feature of all panels in Figure3 is that galaxies divide into two distinct groups, with the top ~20 percent dense environments occupied by galaxies in haloes more massive than ~10!??57! Mo, and the remaining ~80 percent of environments occupied by galaxies in haloes with masses lower than ~10!??A!Mo."," The most noticeable feature of all panels in Figure\ref{fig:NHalo} is that galaxies divide into two distinct groups, with the top $\sim 20$ percent dense environments occupied by galaxies in haloes more massive than $\sim 10^{12.5}\,h^{-1} M_{\odot}$ and the remaining $\sim 80$ percent of environments occupied by galaxies in haloes with masses lower than $\sim 10^{12.5}\,h^{-1} M_{\odot}$."
" This bi-modality arises from the assumed association between galaxies and dark matter haloes required to fit the observed luminosity function and clustering observations, and is explored further below."," This bi-modality arises from the assumed association between galaxies and dark matter haloes required to fit the observed luminosity function and clustering observations, and is explored further below."
" Looking in more detail, the lower 80 percent of rank-ordered densities in Figure 3 shows no trend with halo mass, and as such, the term ‘local environment’ no longer applies."," Looking in more detail, the lower $80$ percent of rank-ordered densities in Figure \ref{fig:NHalo} shows no trend with halo mass, and as such, the term `local environment' no longer applies."
" In terms of a characteristic halo mass for a given environment, this result leaves individual galaxies near clusters indistinguishable from isolated galaxies in voids."," In terms of a characteristic halo mass for a given environment, this result leaves individual galaxies near clusters indistinguishable from isolated galaxies in voids."
" In contrast, the behaviour of the high density-halo mass correlation depends on the neighbour method employed."," In contrast, the behaviour of the high density–halo mass correlation depends on the neighbour method employed."
" In the highest 20 percent environments, low n neighbour searches smooth away any density dependence with halo mass."," In the highest $20$ percent environments, low $n$ neighbour searches smooth away any density dependence with halo mass."
 This can be seen by comparing the farleft panel in Figure 3 (low n) with the far right panel (high m)., This can be seen by comparing the farleft panel in Figure \ref{fig:NHalo} (low $n$ ) with the far right panel (high $n$ ).
" As the number of neighbours used to define environment is increased, galaxies belonging to increasingly massive haloes (which host an increasing number of satellites) will be labelled as increasingly dense."," As the number of neighbours used to define environment is increased, galaxies belonging to increasingly massive haloes (which host an increasing number of satellites) will be labelled as increasingly dense."
" Thus, to more precisely draw out the high density-halo mass environment correlations using nearest neighbour methods, a high n is desirable."," Thus, to more precisely draw out the high density–halo mass environment correlations using nearest neighbour methods, a high $n$ is desirable."
 The first two panels of Figure 3 provide an additional test of the importance of projection effects., The first two panels of Figure \ref{fig:NHalo} provide an additional test of the importance of projection effects.
" Here, the 3rd nearest neighbour count is performed using three dimensional redshift space distances while the 7th nearest neighbour is performed with projected galaxy positions on the two dimensional sky."," Here, the 3rd nearest neighbour count is performed using three dimensional redshift space distances while the 7th nearest neighbour is performed with projected galaxy positions on the two dimensional sky."
 Both methods show the same overall trend with halo mass., Both methods show the same overall trend with halo mass.
" We find that, in general, projecting the galaxy positions simply blurs the edges of the two clouds with the overall shape preserved."," We find that, in general, projecting the galaxy positions simply blurs the edges of the two clouds with the overall shape preserved."
" Another popular neighbour-based method used for measuring environment is Voronoi volumes, as discussed in Section 3.1."," Another popular neighbour-based method used for measuring environment is Voronoi volumes, as discussed in Section \ref{sec:neigh}."
 Figure 4 shows how a Voronoi defined environment estimator also correlates with dark matter halo mass., Figure \ref{fig:vor} shows how a Voronoi defined environment estimator also correlates with dark matter halo mass.
" We see a similar trend to that of the other neighbour- methods, with the overall result close to the 7th nearest projected neighbour method shown in the second panel of Figure 3.."," We see a similar trend to that of the other neighbour-based methods, with the overall result close to the 7th nearest projected neighbour method shown in the second panel of Figure \ref{fig:NHalo}."
 A comparison of Figure 1 with Figures 3 and 4 reveals the origin of the bi-modality., A comparison of Figure \ref{fig:pop} with Figures \ref{fig:NHalo} and \ref{fig:vor} reveals the origin of the bi-modality.
 Galaxies identified to be in the upper 20 percent dense environments tend to be those whose neighbour search stays within the dark matter halo due to a large satellite population., Galaxies identified to be in the upper $20$ percent dense environments tend to be those whose neighbour search stays within the dark matter halo due to a large satellite population.
 Such haloes are almost always more massive than ~10'7°h~1Mo.," Such haloes are almost always more massive than $\sim 10^{12.5}\,h^{-1} M_{\odot}$."
" In contrast, the lower 80 percent density environments are identified by neighbour searches that extend beyond the halo due to a low or zero satellite population of significance."," In contrast, the lower $80$ percent density environments are identified by neighbour searches that extend beyond the halo due to a low or zero satellite population of significance."
" In general, haloes with few satellites almost always have masses smaller than ~1077?57!Ms, and neighbour searches will then tend to probe the inter-halo rather than inter-galaxy separations."," In general, haloes with few satellites almost always have masses smaller than $\sim 10^{12.5}\,h^{-1} M_{\odot}$, and neighbour searches will then tend to probe the inter-halo rather than inter-galaxy separations."
" Many authors have employed fixed apertures to probe the local density around galaxies, as described in Section 3.2.."," Many authors have employed fixed apertures to probe the local density around galaxies, as described in Section \ref{sec:aperture}."
" In a similar vein to Figure 3, Figure 5 shows how various aperture sizes correlate with host dark matter halo mass when a projected fixed aperture is employed with a cut in velocity space around each galaxy."," In a similar vein to Figure \ref{fig:NHalo}, Figure \ref{fig:AHalo} shows how various aperture sizes correlate with host dark matter halo mass when a projected fixed aperture is employed with a cut in velocity space around each galaxy."
" In addition, the central two panels show how the density-halo mass correlation changes if the velocity cut is increased for the same sized aperture."," In addition, the central two panels show how the density–halo mass correlation changes if the velocity cut is increased for the same sized aperture."
" This roughly corresponds to the difference one would expect with data having photometric vs. spectroscopic redshifts,as discussed in Gallazzietal. (2009).."," This roughly corresponds to the difference one would expect with data having photometric vs. spectroscopic redshifts,as discussed in \citet{Gallazzi09}. ."
 The projected fixed aperture technique yields both similar and different trends when compared with the nearest neighbour technique shown in Figure 3.., The projected fixed aperture technique yields both similar and different trends when compared with the nearest neighbour technique shown in Figure \ref{fig:NHalo}. .
 The overall shape, The overall shape
times of 60 seconds.,times of 60 seconds.
 Each series spanned about | hour and was centered on the expected time of the eclipses., Each series spanned about 1 hour and was centered on the expected time of the eclipses.
 The data reduction — bias and flatfield corrections — was performed with IRAF standard routines., The data reduction – bias and flatfield corrections – was performed with IRAF standard routines.
 Ditferential photometry of the images was executed with the aid of the DAOPHOT II package., Differential photometry of the images was executed with the aid of the DAOPHOT II package.
 From the six observed light curves. two were discarded because of the presence of flares that disturbed the measurements of the eclipse timings.," From the six observed light curves, two were discarded because of the presence of flares that disturbed the measurements of the eclipse timings."
 The remaining four eclipse timings were measured individually. as well as combined in a single set folded in phase with the ephemeris from Alcock et al. (1997)).," The remaining four eclipse timings were measured individually, as well as combined in a single set folded in phase with the ephemeris from Alcock et al.\cite{alco}) )."
 The average timing obtained from our data is presented in Table] in addition to the eclipse timings found in the literature., The average timing obtained from our data is presented in \ref{timings} in addition to the eclipse timings found in the literature.
 The ephemeris of Aleock et al. (1997)), The ephemeris of Alcock et al. \cite{alco}) )
" provides a period of P,=0.44267714(6)d; this period was found from observations spanning a total interval of 5884 cycles. with the mean cycle of E,=2942."," provides a period of $P_{1}= 0.442\,677\,14(6)~d$; this period was found from observations spanning a total interval of 5884 cycles, with the mean cycle of $E_{1} = 2942$."
 From our timings we can derive a period of P»=0.44267766(5)d with respect to cycle 5884: this corresponds to the mean cycle E»=9915.," From our timings we can derive a period of $P_{2} = 0.442\,677\,66(5)~d$ with respect to cycle 5884; this corresponds to the mean cycle $E_{2} = 9915$."
" The difference between P» and P, is 9 to 10 times the uncertainty in the determination of those periods.", The difference between $P_{2}$ and $P_{1}$ is 9 to 10 times the uncertainty in the determination of those periods.
" From the difference between Ps and P, and the elapsed time P(E»—Ej) we obtain P=+1.7(40.3)x 107. the uncertainty of which was estimated with the usual error propagation method from the errors of Pi and P»."," From the difference between $P_{2}$ and $P_{1}$ and the elapsed time $P.(E_2-E_1)$ we obtain $\dot{P} = +1.7(\pm 0.3) \times 10^{-10}$ , the uncertainty of which was estimated with the usual error propagation method from the errors of $P_{1}$ and $P_{2}$."
 As expected. the new timing detemination allowed the measurement of the orbital period change: the period increases with a time-scale of P/P=+7.2(41.3)x10° years.," As expected, the new timing detemination allowed the measurement of the orbital period change: the period increases with a time-scale of $P/\dot{P} = +7.2(\pm1.3) \times 10^{6}$ years."
 The first conclusion is that. contrary to the canonical view that considers CBSS as having massive (1-2 Ma) secondary stars. CAL 87 clearly has a mass ratio which is typical for cataclysmic variables. te.. M»«Mj.," The first conclusion is that, contrary to the canonical view that considers CBSS as having massive (1–2 $M_{\sun}$ ) secondary stars, CAL 87 clearly has a mass ratio which is typical for cataclysmic variables, i.e., $M_2<M_1$."
 This has been anticipated by Cowley et al. (1990... 1998))," This has been anticipated by Cowley et al. \cite{cow}, \cite{cow98}) )"
" who found. given the small amplitude of the radial velocity of emission lines. that M»=0.43Mc and M,>4M."," who found, given the small amplitude of the radial velocity of emission lines, that $M_2 = 0.4 M_{\sun}$ and $M_1 > 4 M_{\sun}$."
 For this reason those authors explored the idea that such a massive primarycould be a black hole., For this reason those authors explored the idea that such a massive primarycould be a black hole.
 We now know that the mass accreting star is probably a white dwarf: its mass. derived. from fitting the Iuminosity-temperature diagram with theoretical calculations. is (Starrfield et al. 20049).," We now know that the mass accreting star is probably a white dwarf: its mass, derived from fitting the luminosity-temperature diagram with theoretical calculations, is (Starrfield et al. \cite{starr}) )."
 The observed supersoft X-ray luminosity. Ly=4x10°° erg/s (Starrfield et al. 20049).," The observed supersoft X-ray luminosity, $L_X = 4 \times 10^{36}$ erg/s (Starrfield et al. \cite{starr}) ),"
" implies an aceretion rate of M,=107Mevr!.", implies an accretion rate of $\dot{M}_{acc}= 10^{-8} M_{\sun}~yr^{-1}$.
 This is additional information which can be used to solve the following equations (van Teeseling King 1998:: see also eqs., This is additional information which can be used to solve the following equations (van Teeseling King \cite{tees}; see also eqs.
 3. 4 and 5 in Steiner et al. 2006)):," 3, 4 and 5 in Steiner et al. \cite{stei06}) ):"
 and where Here f» is q2-the 3g)ratio of the specific angular momentum of the wind to that of the secondary star., and where Here $\beta_2$ is the ratio of the specific angular momentum of the wind to that of the secondary star.
 We have also assumed that M»«0.6 (see van Teeseling King 1998. for a detailed discussion)., We have also assumed that $M_2 < 0.6$ (see van Teeseling King \cite{tees} for a detailed discussion).
 The right-hand side of equations (1)) and (2)) are measured quantities. while from the left-hand side two independent parameters can be determined: g=0.25 and f»=0.89.," The right-hand side of equations \ref{eq:um}) ) and \ref{eq:dois}) ) are measured quantities, while from the left-hand side two independent parameters can be determined: $q = 0.25$ and $\beta_2 = 0.89$."
 This implies that the mass of the secondary is M»=0.34Moe., This implies that the mass of the secondary is $M_2 = 0.34 M_{\sun}$.
 One should explore the range of parameters that could be accommodated with the observations., One should explore the range of parameters that could be accommodated with the observations.
" The observed supersoft X-ray luminosity means that M, cannot be smaller than 1073Me, vr!: therefore the minimum values for B»and q are 0.89 and 0.25. respectively."," The observed supersoft X-ray luminosity means that $\dot{M}_{acc}$ cannot be smaller than $10^{-8} M_{\sun}~yr^{-1}$ ; therefore the minimum values for $\beta_2$and $q$ are 0.89 and 0.25, respectively."
 A larger value for M... could be possible if part of the X-rays are absorbed. for instance. due to the high inclination.," A larger value for $\dot{M}_{acc}$ could be possible if part of the X-rays are absorbed, for instance, due to the high inclination."
 However. equation (3)) dictates thatq« 2/3.," However, equation \ref{eq:tres}) ) dictates that$q<2/3$ ."
" For this reason Mi,«1.9x107Mg, vr!. but M» could be as high as M»=0.9 Mo."," For this reason $\dot{M}_{acc} < 1.9 \times 10^{-8} M_{\sun}~yr^{-1}$ , but $M_{2}$ could be as high as $M_{2}=0.9~M_{\sun}$ ."
 For all the range of, For all the range of
"But, for the limiting case in which outward migration can only occur in an isothermal disk, we still expect the Grand Tack to happen because in the last stages of the disk's lifetime it is necessarily optically thin and, as seen in Sect.","But, for the limiting case in which outward migration can only occur in an isothermal disk, we still expect the Grand Tack to happen because in the last stages of the disk's lifetime it is necessarily optically thin and, as seen in Sect."
" 3, this leads inevitably to outward migration."," 3, this leads inevitably to outward migration."
" Our results indicate that Jupiter and Saturn probably underwent a two-phase, inward-then-outward migration."," Our results indicate that Jupiter and Saturn probably underwent a two-phase, inward-then-outward migration."
" In our simulations, Jupiter and Saturn start as 10 Mg cores and type I migrate; inward for isothermal disks, inward or outward for radiative disks."," In our simulations, Jupiter and Saturn start as $10$ $\mearth$ cores and type I migrate; inward for isothermal disks, inward or outward for radiative disks."
 In most cases the two cores become locked in 3:2 mean motion resonance (MMR)., In most cases the two cores become locked in 3:2 mean motion resonance (MMR).
" At this point or after a delay of 500-1000 orbits, we allowed Jupiter to start accreting gas from the disk."," At this point or after a delay of 500-1000 orbits, we allowed Jupiter to start accreting gas from the disk."
" When Jupiter reaches the gap-opening mass, it undergoes a phase of rapid inward migration as it clears out its gap (sometimes called type III migration; Masset Papaloizou 2003) then settles into standard, type II migration."," When Jupiter reaches the gap-opening mass, it undergoes a phase of rapid inward migration as it clears out its gap (sometimes called type III migration; Masset Papaloizou 2003) then settles into standard, type II migration."
 Inward migration continues until Saturn accretes enough gas to reach the gap-opening mass itself., Inward migration continues until Saturn accretes enough gas to reach the gap-opening mass itself.
" At this point, Saturn's inward migration accelerates and is again trapped in the 3:2 MMR with Jupiter."," At this point, Saturn's inward migration accelerates and is again trapped in the 3:2 MMR with Jupiter."
 Outward migration of both giant planets is then triggered via the mechanism of Masset Snellgrove (2001)., Outward migration of both giant planets is then triggered via the mechanism of Masset Snellgrove (2001).
" Outward migration stops when either a) the disk dissipates (as in Sect. 3.2.5)),"," Outward migration stops when either a) the disk dissipates (as in Sect. \ref{sec:disp}) ),"
" b) Saturn reaches the outer edge of the disk, or, c) if the disk is flared, the giant planets drop below the local gap-opening mass (e.g., Crida et al."," b) Saturn reaches the outer edge of the disk, or, c) if the disk is flared, the giant planets drop below the local gap-opening mass (e.g., Crida et al."
 2009)., 2009).
 An additional stopping — or at least slowing — mechanism exists if the planets are unable to maintain a well-aligned resonant lock during migration., An additional stopping – or at least slowing – mechanism exists if the planets are unable to maintain a well-aligned resonant lock during migration.
" For example, in simulation I1 Jupiter and Saturn's rate of outward migration slowed significantly when one resonant angle transitioned from libration to circulation (see Figs."," For example, in simulation I1 Jupiter and Saturn's rate of outward migration slowed significantly when one resonant angle transitioned from libration to circulation (see Figs."
 1 and 3)., 1 and 3).
" Here, this appears to be due to a positive corotation torque exerted on Saturn by gas that polluted Jupiter and Saturn's common gap as the distance between the two planets increased during outward migration."," Here, this appears to be due to a positive corotation torque exerted on Saturn by gas that polluted Jupiter and Saturn's common gap as the distance between the two planets increased during outward migration."
 On longer timescales it is unclear if this mechanism would continue to slow down and eventually stop the outward migration., On longer timescales it is unclear if this mechanism would continue to slow down and eventually stop the outward migration.
" In isothermal disks, the two phase migration of Jupiter and Saturn holds for almost the full range of parameters that we tested (Sect."," In isothermal disks, the two phase migration of Jupiter and Saturn holds for almost the full range of parameters that we tested (Sect."
 3)., 3).
 The only situation for which this result does not hold is if the gaseous Solar Nebula is relatively thick (h=A/r 0.05)., The only situation for which this result does not hold is if the gaseous Solar Nebula is relatively thick $h = H/r \gtrsim 0.05$ ).
 Both the disk’s surface density profile and the value for the disk aspect ratio had an effect on the migration rate: disks with either shallower profiles or lower values of h result in faster migration., Both the disk's surface density profile and the value for the disk aspect ratio had an effect on the migration rate: disks with either shallower profiles or lower values of $h$ result in faster migration.
" Changing the disk's viscosity had little effect on the outcome, although we only tested a very small range."," Changing the disk's viscosity had little effect on the outcome, although we only tested a very small range."
" In one simulation (I4; Sect. 3.2.1)),"," In one simulation (I4; Sect. \ref{sec:varyxj}) ),"
" Saturn's core was pushed past the 2:1 MMR with Jupiter leading to a significant eccentricity increase for both planets before the resonance was crossed, Saturn was trapped in the 2:3 MMR and both planets migrated outward."," Saturn's core was pushed past the 2:1 MMR with Jupiter leading to a significant eccentricity increase for both planets before the resonance was crossed, Saturn was trapped in the 2:3 MMR and both planets migrated outward."
" This dynamic phase of resonance crossing and eccentricity excitation is likely to be quite sensitive to the detailed properties of the disk (e.g., the scale height and viscosity) that determine the gap profile."," This dynamic phase of resonance crossing and eccentricity excitation is likely to be quite sensitive to the detailed properties of the disk (e.g., the scale height and viscosity) that determine the gap profile."
 We performed two simulations in radiative disks with mixed outcomes (Sect., We performed two simulations in radiative disks with mixed outcomes (Sect.
 4)., 4).
" In the first case, Jupiter and Saturn started accreting together and so stayed relatively close to each other."," In the first case, Jupiter and Saturn started accreting together and so stayed relatively close to each other."
 The planets became locked in the 3:2 MMR and migrated outward even faster than in isothermal simulations, The planets became locked in the 3:2 MMR and migrated outward even faster than in isothermal simulations
"spatially flat ACDM models characterized bv the matter density. parameter Q,,. vacuum enerev densitv parameter O4 and Hubble constant .=0.75. and will calculate lensine probabilities with image separations greater than A according to GNEW halo profile models with varying a.","spatially flat $\Lambda$ CDM models characterized by the matter density parameter $\Omega_{\mathrm
m}$, vacuum energy density parameter $\Omega_{\Lambda}$ and Hubble constant $h=0.75$, and will calculate lensing probabilities with image separations greater than $\Delta\theta$ according to GNFW halo profile models with varying $\alpha$."
" In the definition of cross-section or lensing probability. we will just consider the criterion Nd. and neglect. another one q, (the brightness ratio of the outermost two images) which is (he ratio of (he corresponding absolute values of the maenilications (?).."," In the definition of cross-section or lensing probability, we will just consider the criterion $\Delta\theta$, and neglect another one $q_r$ (the brightness ratio of the outermost two images) which is the ratio of the corresponding absolute values of the magnifications \citep{1992grle.book.....S}."
" In order to investigate (he effect of central black holes or bulges on lensing probability. introduced g, into the calculation for lensing cross-section."," In order to investigate the effect of central black holes or bulges on lensing probability, \cite{2003ApJ...587L..55C,2003A&A...397..415C} introduced $q_r$ into the calculation for lensing cross-section."
" Due to the existence of central black holes or galactic bulges. y,,. becomes extremely large when |r| approaches zero."," Due to the existence of central black holes or galactic bulges, $y_{cr}$ becomes extremely large when $|x|$ approaches zero."
" Thus Yer CoN be determined by the consideration of ¢, together with AG.", Thus $y_{cr}$ can be determined by the consideration of $q_r$ together with $\Delta\theta$.
 ILowever lor GNEW halo nodels in the absence of central black holes or galactic bulges. we can see in Fig.l that the ensing equation curves are so smooth that we do not need (o define cross-section by ο.," However for GNFW halo models in the absence of central black holes or galactic bulges, we can see in \ref{fig:le} that the lensing equation curves are so smooth that we do not need to define cross-section by $q_r$."
" Our numerical results are shown in Fig.2 together with the observational ones. which or 6<AO«15"" [rom JVAS/CLASS takes the form of an upper limit."," Our numerical results are shown in \ref{fig:lp} together with the observational ones, which for $6^{''}\leq\Delta\theta\leq
15^{''}$ from JVAS/CLASS takes the form of an upper limit."
 Based on the ugh-vresolution simulation results (?).. we adopt a compound GNEW model for dark halos with three populations as described above.," Based on the high-resolution simulation results \citep{2000ApJ...529L..69J}, we adopt a compound GNFW model for dark halos with three populations as described above."
" For the cosmological models. we first choose the jew result from the Wilkinson Microwave Anisotropy Probe (WMADP): O4,=0.27. σε=0.84 (2?) and M4=LOMAS... the numerical result of which is depicted by a thin solid line in Fig.2.."," For the cosmological models, we first choose the new result from the Wilkinson Microwave Anisotropy Probe (WMAP): $\Omega_{\mathrm
m}=0.27$, $\sigma_8=0.84$ \citep{2003ApJS..148....1B,2003ApJS..148..175S} and $M_{c1}=10^{13}M_{\sun}$, the numerical result of which is depicted by a thin solid line in \ref{fig:lp}. ."
 lt is clear (hat the predicted lensing probabilities are much lower than the observed values., It is clear that the predicted lensing probabilities are much lower than the observed values.
" Previousworks (727) showed that lensing probability is sensitive to O,,. oy and Ma."," Previousworks \citep{2003ApJ...587L..55C,2003A&A...397..415C} showed that lensing probability is sensitive to $\Omega_{\mathrm m}$, $\sigma_8$ and $M_{c1}$."
" More specifically. it increases with 9,4. e; and M4."," More specifically, it increases with $\Omega_{\mathrm m}$, $\sigma_8$ and $M_{c1}$."
 Ir order to improve the prediction. we take Ou. ox and M4 to be the higher values of 0.4. 1.2 and 5x10M. respectively.," In order to improve the prediction, we take $\Omega_{\mathrm m}$, $\sigma_8$ and $M_{c1}$ to be the higher values of 0.4, 1.2 and $5\times10^{13}M_{\sun}$ respectively."
 As shown bv the dashed line in Fig.2.. the derived lensine probabilities are still inconsistent with the observations although thev have increased erveatly compared to the former moclel.," As shown by the dashed line in \ref{fig:lp}, the derived lensing probabilities are still inconsistent with the observations although they have increased greatly compared to the former model."
 Thus we take (he upper limit of 1.6 lor e. as used in literature to date without changing the other parameters.," Thus we take the upper limit of 1.6 for $\sigma_8$, as used in literature to date without changing the other parameters."
 The results are shown with a doted line in Fie.2.., The results are shown with a doted line in \ref{fig:lp}.
 There is almost no improvement in the small angle separations compared with the case of σε=1.2. apart from the increase of the predicted probabilities on large angles.," There is almost no improvement in the small angle separations compared with the case of $\sigma_8=1.2$, apart from the increase of the predicted probabilities on large angles."
 On the other hancl. the lensing probabilities for both of cases σε=1.2 and σς=1.6 lie above the observed upper limit for large angle separations.," On the other hand, the lensing probabilities for both of cases $\sigma_8=1.2$ and $\sigma_8=1.6$ lie above the observed upper limit for large angle separations."
 We have generalized the censity profiles from high-resolution simulations of dark matter halos accordinge to GNEW modelswith varvinge a. and developed a compound model incorporatinge three populations of matter halos.," We have generalized the density profiles from high-resolution simulations of dark matter halos according to GNFW modelswith varying $\alpha$ , and developed a compound model incorporating three populations of matter halos."
 We choose three tvpical cosmological models: Model A.," We choose three typical cosmological models: Model A,"
approaches to dealing with realistic coagulation kernels that employ the moments method under (he assumption that the form of the mass distribution. f(r.) is à powerlaw.,"approaches to dealing with realistic coagulation kernels that employ the moments method under the assumption that the form of the mass distribution $f(m,t)$ is a powerlaw."
 Such an assumption is not entirely unfounded., Such an assumption is not entirely unfounded.
 A number of detailed models by Weidenschilling2004) have shown that powerlaw size distributions result. which have nearly constant mass per decade radius to an upper mit (ΗΕ) which grows with time until the frustration limit of around a meter is reached. and (under turbulent conditions. at least) erowth stalls (Cuzzi&Weidenschilling2006).," A number of detailed models by \citet{wei97,wei00,wei04}
 have shown that powerlaw size distributions result, which have nearly constant mass per decade radius to an upper limit $m_L(t)$ which grows with time until the frustration limit of around a meter is reached, and (under turbulent conditions, at least) growth stalls \citep{cuz06}."
. Similar trends are found by Dullemoncl&Dominik(2005) in which dillerent assumptions about collisional ejecta are made., Similar trends are found by \citet{dul05} in which different assumptions about collisional ejecta are made.
 These authors do find minor fluctuations in the distribution. but if one is primarily interested in eeneral properties of the distribution. such as how the largest particle size ehanges with time. and not the fine structure of (he mass distribution itself. (he approach is quite advantageous.," These authors do find minor fluctuations in the distribution, but if one is primarily interested in general properties of the distribution, such as how the largest particle size changes with time, and not the fine structure of the mass distribution itself, the approach is quite advantageous."
 There are reasons to believe that a powerlaw is a natural end-state. especially those with equal mass per decade. because they have sell-preserving properties for the collision kernels ol interest (Cuzzi&Weidenschilling 2006)..," There are reasons to believe that a powerlaw is a natural end-state, especially those with equal mass per decade, because they have self-preserving properties for the collision kernels of interest \citep{cuz06}. ."
" The significance of the upper mass cutoff me, depends on the assumed slope of the powerlaw.", The significance of the upper mass cutoff $m_L$ depends on the assumed slope of the powerlaw.
 In all real distributions. (here will be a rapidly decreasing abundance of particles for masses exceeding some threshold. even though the abundance may not drop immnediately lo zero as in our assunied model.," In all real distributions, there will be a rapidly decreasing abundance of particles for masses exceeding some threshold, even though the abundance may not drop immediately to zero as in our assumed model."
 The rare. extremely large particles might be of interest for sole applications. but our locus will be the particle size carrving most of (he area. or 1osl of the mass (the first or second moments).," The rare, extremely large particles might be of interest for some applications, but our focus will be the particle size carrying most of the area, or most of the mass (the first or second moments)."
 For powerlaw distributions (i)xm© with q2. my is ilself the mass of the particle earrving most of the mass and is independent of the selection of a lower particle size cutoff.," For powerlaw distributions $f(m) \propto m^{-q}$ with $q < 2$, $m_L$ is itself the mass of the particle carrying most of the mass and is independent of the selection of a lower particle size cutoff."
 For q=2. the distribution contains equal nass per decade. and [for q> 2. most of the mass is in the small particles and (he value of my depends on the lower particle size chosen.," For $q=2$, the distribution contains equal mass per decade, and for $q>2$ , most of the mass is in the small particles and the value of $m_L$ depends on the lower particle size chosen."
 Most realistic distributions. and (hose most commonly treated in the literature. have gq<2: here we assume gq=11/6. a widely used Iragmentation powerlaw.," Most realistic distributions, and those most commonly treated in the literature, have $q<2$; here we assume $q=11/6$, a widely used fragmentation powerlaw."
 In this case. a strict cutoff al mp represents a well-defined distribution with easilv-understood moments where most of the mass is at my and the area is nearly equally distributed per decade with a mass dependence mE.," In this case, a strict cutoff at $m_L$ represents a well-defined distribution with easily-understood moments where most of the mass is at $m_L$ and the area is nearly equally distributed per decade with a mass dependence $m^{-1/6}$."
 Although we do not include either imperfect sticking or fragmentation al (his stage in the model. we believe that the moment equations as expressed in Eq. (," Although we do not include either imperfect sticking or fragmentation at this stage in the model, we believe that the moment equations as expressed in Eq. ("
5) will remain valid up to a Tragmentation barrier”. which may be defined as that size for which the (vpical disruption energv of a particle is on (he same order as the enerev of identical colliding particles.,"5) will remain valid up to a ""fragmentation barrier"", which may be defined as that size for which the typical disruption energy of a particle is on the same order as the energy of identical colliding particles."
 The fragmentation barrier will also depend on ones choice of nebula parameters., The fragmentation barrier will also depend on one's choice of nebula parameters.
 This treatment (up to the fragmentation size) is consistent wilh recent work by (their Fie., This treatment (up to the fragmentation size) is consistent with recent work by \citet{bra08} (their Fig.
 13) which shows a constant powerlaw mass distribution up(ο a cutloll size which thenfalls olf abruptly., 13) which shows a constant powerlaw mass distribution upto a cutoff size which thenfalls off abruptly.
 This “knee” in the distribution represents that efficient fragmentation size., This “knee” in the distribution represents that efficient fragmentation size.
The amount of IR-excess in a disk is directly related to its geometry (Kenyon&Hartmann1987;Meeusetal.2001;Dullemondet 2001).,"The amount of IR-excess in a disk is directly related to its geometry \citep{KE87,GM01,DU01}."
". Specifically using the IRS spectra, disk geometry can be inferred from the flux ratio between 30 and 13 jum (F3o/F13,Brownetal.2007;Oliveiraetal.2010;Merín 2010))."," Specifically using the IRS spectra, disk geometry can be inferred from the flux ratio between 30 and 13 $\mu$ m $F_{30}/F_{13}$,\citealt{JB07,OL10,ME10}) )."
" A flared geometry S 5), with considerable IR excess and (1.5small dust,Ps3o/Fia allows the uppermost dust layers to intercept stellar light at both the inner and outer disk."," A flared geometry $1.5
\lesssim F_{30}/F_{13} \lesssim 5$ ), with considerable IR excess and small dust, allows the uppermost dust layers to intercept stellar light at both the inner and outer disk."
" For flat disks (Fso/ 1.5) with little IR excess, only the inner disk can easilyΕις S;intercept the stellar radiation as the outer disk is shadowed."," For flat disks $F_{30}/F_{13}
\lesssim 1.5$ ) with little IR excess, only the inner disk can easily intercept the stellar radiation as the outer disk is shadowed."
" Moreover, cold or transitional disks are interesting objects that present inner dust gaps or holes, producing a region with little or no near-IR excess (5S;Fso/Fia 15)."," Moreover, cold or transitional disks are interesting objects that present inner dust gaps or holes, producing a region with little or no near-IR excess $5 \lesssim F_{30}/F_{13} \lesssim 15$ )."
 It is interesting to explore the effect of disk geometry on both the mean mass-average grain sizes and crystallinity fractions of the disks studied., It is interesting to explore the effect of disk geometry on both the mean mass-average grain sizes and crystallinity fractions of the disks studied.
 Figure 4 shows F'39/F13 as a proxy for disk geometry compared with the mean mass-averaged grain sizes and crystallinity fractions for both components and all regions studied here., Figure \ref{f_f30} shows $F_{30}/F_{13}$ as a proxy for disk geometry compared with the mean mass-averaged grain sizes and crystallinity fractions for both components and all regions studied here.
" No preferential grain size coefficient 7 — -0.14, P — 0.02, and 7 — 0.07, P = 0.33 (correlationfor warm and cold components, respectively) nor crystallinity fraction (r — 0.09, P — 0.10 for the warm, and 7 — -0.19, P — 0.01 for the cold component) is apparent for any given disk geometry."," No preferential grain size (correlation coefficient $\tau$ = -0.14, $P$ = 0.02, and $\tau$ = 0.07, $P$ = 0.33 for warm and cold components, respectively) nor crystallinity fraction $\tau$ = 0.09, $P$ = 0.10 for the warm, and $\tau$ = -0.19, $P$ = 0.01 for the cold component) is apparent for any given disk geometry."
" Similar scatter plots result for the mean mass-average grains sizes for only amorphous (τς -0.12, P = 0.08 for the warm, and 7 = 0.13, P = 0.11 for the cold component), or only crystalline grains = 0.08, P = 0.17 for the warm, and 7 — -0.13, P — 0.10 (7for the cold component)."," Similar scatter plots result for the mean mass-average grains sizes for only amorphous $\tau$ = -0.12, $P$ = 0.08 for the warm, and $\tau$ = 0.13, $P$ = 0.11 for the cold component), or only crystalline grains $\tau$ = 0.08, $P$ = 0.17 for the warm, and $\tau$ = -0.13, $P$ = 0.10 for the cold component)."
" Furthermore, no clear separation is seen between the different regions studied."," Furthermore, no clear separation is seen between the different regions studied."
" The statistically relevant samples in Serpens and Taurus define a locus where the majority of the objects is located in each plot, which is followed by the lower number statistics for older regions."," The statistically relevant samples in Serpens and Taurus define a locus where the majority of the objects is located in each plot, which is followed by the lower number statistics for older regions."
" Figure 4 therefore shows not only that grain size and crystallinity fraction are not a function of disk geometry, but also that younger and older regions show similar distributions of those two parameters."," Figure \ref{f_f30} therefore shows not only that grain size and crystallinity fraction are not a function of disk geometry, but also that younger and older regions show similar distributions of those two parameters."
" The crystallinity fractions derived from the warm and cold components (Cwarm and σοια, respectively) for Serpens and Taurus are show in Figure 5.."," The crystallinity fractions derived from the warm and cold components $C_{\rm Warm}$ and $C_{\rm Cold}$, respectively) for Serpens and Taurus are show in Figure \ref{f_cryst}."
" No strong trend of warm and cold crystallinity fractions increasing together is seen (r = 0.10, with P = 0.10 for the entire sample)."," No strong trend of warm and cold crystallinity fractions increasing together is seen $\tau$ = 0.10, with $P$ = 0.10 for the entire sample)."
" This fact implies that, if an unique process is responsible for the crystallization of dust at all radii, this process is not occurring at the same rate in the innermost regions as further out in the disk."," This fact implies that, if an unique process is responsible for the crystallization of dust at all radii, this process is not occurring at the same rate in the innermost regions as further out in the disk."
" This is opposite to the conclusion of Watsonetal.(2009),, who derive a correlation between inner and outer disk crystallinity from the simultaneous presence of the 11.3 and 33 um features."," This is opposite to the conclusion of \citet{WA09}, who derive a correlation between inner and outer disk crystallinity from the simultaneous presence of the 11.3 and 33 $\mu$ m features."
" The opacities of the crystalline species are more complex than those two features alone, making the analysis here more complete than that of Watson"," The opacities of the crystalline species are more complex than those two features alone, making the analysis here more complete than that of \citet{WA09}. ."
 Our finding that the fraction of crystalline material in disk surfaces varies with radius can constrain some of the mechanisms for formation and distribution of crystals., Our finding that the fraction of crystalline material in disk surfaces varies with radius can constrain some of the mechanisms for formation and distribution of crystals.
 A wider spread in crystallinity fraction is observed, A wider spread in crystallinity fraction is observed
"Since the relativistic line and disk reflection models are necessarily separate when ""pexrav"" is employed. the equivalent width of the relativistic line can be measured directly.","Since the relativistic line and disk reflection models are necessarily separate when “pexrav” is employed, the equivalent width of the relativistic line can be measured directly."
 We find a line equivalent width of W=130+10 eV. Each fit made with pexrav found Fairall 9 to be consistent with low or moderate spin values., We find a line equivalent width of $W = 130\pm 10$ eV. Each fit made with pexrav found Fairall 9 to be consistent with low or moderate spin values.
 Abundances of 0.5 and 1.0 gave spin values of q—0.077 and a=0.1751. respectively.," Abundances of 0.5 and 1.0 gave spin values of $a = 0.0^{+0.2}$ and $ a =
0.1_{-0.1}^{+0.5}$, respectively."
 In all cases. maximal spin is excluded at more than the 5c level of confidence.," In all cases, maximal spin is excluded at more than the $5\sigma$ level of confidence."
 Since the convolution model has the inner disk inclination as a variable parameter. while pexrav uses the cosine of that angle. it was not possible to link these two parameters directly.," Since the convolution model has the inner disk inclination as a variable parameter, while pexrav uses the cosine of that angle, it was not possible to link these two parameters directly."
 The results discussed above are based on an inclination of 40 degrees., The results discussed above are based on an inclination of 40 degrees.
 This value was selected after fitting the model with several different inclination values and tracing the evolution of the goodness-of-fit statistic. and it is 1n broad agreement with the inclination found using other models (see below).," This value was selected after fitting the model with several different inclination values and tracing the evolution of the goodness-of-fit statistic, and it is in broad agreement with the inclination found using other models (see below)."
 The results we obtained with this model are not entirely self-consistent., The results we obtained with this model are not entirely self-consistent.
 A relativistic line centroid energy of 6.7050! keV is measured. consistent with He-like Fe XXV.," A relativistic line centroid energy of $6.70^{+0.01}_{-0.03}$ keV is measured, consistent with He-like Fe XXV."
 This is at odds with the assumption of a completely neutral aceretion disk., This is at odds with the assumption of a completely neutral accretion disk.
 Moreover. a reflection fraction of 2.0 1s required in all fits.," Moreover, a reflection fraction of 2.0 is required in all fits."
 Yet a reflection fraction of 70.7 is suggested by the equivalent width of the relativistic line (George Fabian 1991)., Yet a reflection fraction of $\simeq0.7$ is suggested by the equivalent width of the relativistic line (George Fabian 1991).
" Our best fits to the data were obtained using ""reflionx (Ross Fabian 2005).", Our best fits to the data were obtained using “reflionx” (Ross Fabian 2005).
 This model includes an Fe K emission line. à broad range of ionization. species. and allows the iron abundance to be a free parameter.," This model includes an Fe K emission line, a broad range of ionization species, and allows the iron abundance to be a free parameter."
 Unlike pexrav. it is an angle-averaged model: the inclination angle is not a variable paramter in spectral fits.," Unlike pexrav, it is an angle-averaged model; the inclination angle is not a variable paramter in spectral fits."
 Owing to the fact that reflionx includes low energy emission lines that can be blurred into a pseudocontinuum that could be the origin of the soft excess (e.g. Crummy et 22006). we did not include a disk blackbody when fitting with reflionx.," Owing to the fact that reflionx includes low energy emission lines that can be blurred into a pseudocontinuum that could be the origin of the soft excess (e.g. Crummy et 2006), we did not include a disk blackbody when fitting with reflionx."
 It should be noted that reflionx. requires a photon power-law index as an input but requires a separate power-law component to fit the continuum., It should be noted that reflionx requires a photon power-law index as an input but requires a separate power-law component to fit the continuum.
 Acordingly. we linked the power-law index between the two components.," Acordingly, we linked the power-law index between the two components."
 Acceptable fits to Fairall 9 could not be obtained when the Fe abundance and line emissivity parameters were frozen at their nominal values., Acceptable fits to Fairall 9 could not be obtained when the Fe abundance and line emissivity parameters were frozen at their nominal values.
 These parameters were therefore allowed to vary., These parameters were therefore allowed to vary.
 With 5970 degrees of freedom. a blurred reflionx model gave à good fit: 4=6498.6 (see Table | and Figure 4).," With 5970 degrees of freedom, a blurred reflionx model gave a good fit: $\chi^{2} = 6498.6$ (see Table 1 and Figure 4)."
 This is significantly better than the fits achieved with the pexrav., This is significantly better than the fits achieved with the pexrav.
" Whereas fits with “pexrav” required an ionized line despite the assumption of neutral reflection. ""reflionx"" returns more self-consistent and reasonable parameter values."," Whereas fits with “pexrav” required an ionized line despite the assumption of neutral reflection, “reflionx” returns more self-consistent and reasonable parameter values."
 A steep line emissivity is required (a maximum of q=5 was fixed as per the case of light bending near to a spinning black hole Miniutti et 22003)., A steep line emissivity is required (a maximum of $q=5$ was fixed as per the case of light bending near to a spinning black hole Miniutti et 2003).
 Using reflionx. we measure a spin of a=0.603:0.07.," Using reflionx, we measure a spin of $a = 0.60\pm 0.07$."
 A maximal spin ts ruled out at the 10c level of confidence. and zero spin is ruled out at more than the 7c level of confidence (see Figure 5).," A maximal spin is ruled out at the $10\sigma$ level of confidence, and zero spin is ruled out at more than the $\sigma$ level of confidence (see Figure 5)."
 Owing to the fact that CCDs are made of Si. effective area curves often change abruptly in the Si band.," Owing to the fact that CCDs are made of Si, effective area curves often change abruptly in the Si band."
 This makes it difficult to calibrate detector responses in the Si range., This makes it difficult to calibrate detector responses in the Si range.
 A formally acceptable fit with reflionx is not found only due to lingering difficulties in the calibration of the XIS response in the Si band., A formally acceptable fit with reflionx is not found only due to lingering difficulties in the calibration of the XIS response in the Si band.
 In Figure |. clear residuals are seen in this narrow band that differ between the XIS cameras.," In Figure 1, clear residuals are seen in this narrow band that differ between the XIS cameras."
 The residuals do not affect the broad-band fit parameters apart from the statistic., The residuals do not affect the broad-band fit parameters apart from the statistic.
" To understand the influence of the soft X-ray band on the spin constraint made with ""reflionx. we ignored the spectra below 2 keV and performed new error scans on the spin parameter."," To understand the influence of the soft X-ray band on the spin constraint made with “reflionx”, we ignored the spectra below 2 keV and performed new error scans on the spin parameter."
 These fits achieve à significantly worse spin constraint: «20.5751.," These fits achieve a significantly worse spin constraint: $a =
0.5^{+0.1}_{-0.3}$."
 Zero spin is only excluded at the level of confidence: however. maximal spin is excluded at more than the Sc Though it is common to attribute the spectral features we have observed to X-ray reflection. from the inner disk. it is important to rigorously rule out alternatives.," Zero spin is only excluded at the level of confidence; however, maximal spin is excluded at more than the $5\sigma$ Though it is common to attribute the spectral features we have observed to X-ray reflection from the inner disk, it is important to rigorously rule out alternatives."
" This is especially important in cases like Fairall 9: although the continuum is arguably ""simpler"" than that of Seyfert | AGN that have an X-ray warm absorber (see. e.g. Blustin et 22005). the signal to noise ratio in the Fe K band is lower than in cases like MCG- (e.g. Miniutti et 22007)."," This is especially important in cases like Fairall 9: although the continuum is arguably “simpler” than that of Seyfert 1 AGN that have an X-ray warm absorber (see, e.g. Blustin et 2005), the signal to noise ratio in the Fe K band is lower than in cases like MCG-6-30-15 (e.g. Miniutti et 2007)."
 We conducted additional investigations to evaluate the robustness of the disk relection spectrum and relativistic line interpretation., We conducted additional investigations to evaluate the robustness of the disk relection spectrum and relativistic line interpretation.
 The spectra shown in Figures |. 2. and 3 — and evidence for a relativistic disk line — could be biased by the influence of a disk reflection spectrum.," The spectra shown in Figures 1, 2, and 3 – and evidence for a relativistic disk line – could be biased by the influence of a disk reflection spectrum."
" We therefore replaced the power-law model used in those figures with a ""pexrav"" reflection model.", We therefore replaced the power-law model used in those figures with a “pexrav” reflection model.
 The reflection fraction was set to 0.7. which corresponds to the equivalent width of the neutral Fe K line at 6.4 keV (R=EW/180 eV: George Fabian 1991).," The reflection fraction was set to 0.7, which corresponds to the equivalent width of the neutral Fe K line at 6.4 keV ${\rm R} = {\rm EW}/180$ eV; George Fabian 1991)."
 Note that this is itself conservative. since the reflection edge is a sharp neutral edge. not one that is broadened by Compton scattering. and the reflection spectrum was not blurred.," Note that this is itself conservative, since the reflection edge is a sharp neutral edge, not one that is broadened by Compton scattering, and the reflection spectrum was not blurred."
 Two neutral Gaussians corresponding to neutral Fe Ko. and K./ lines were added as before., Two neutral Gaussians corresponding to neutral Fe $\alpha$ and $\beta$ lines were added as before.
 A poor fit is achieved with this model (47/»= 6747.4/5975): broad residuals are still visible in the spectrum (see Figure 6)., A poor fit is achieved with this model $\chi^{2}/\nu = 6747.4/5975$ ); broad residuals are still visible in the spectrum (see Figure 6).
 Next. two additional narrow Gaussians corresponding to like Fe XXV and H-like Fe XXVI were included in the model.," Next, two additional narrow Gaussians corresponding to He-like Fe XXV and H-like Fe XXVI were included in the model."
 These lines were allowed to vary freely in flux., These lines were allowed to vary freely in flux.
 This is a very conservative measure: narrow lines corresponding to Fe XXV, This is a very conservative measure: narrow lines corresponding to Fe XXV
where is the Fourier triusforii of the window function Were). aud jy is a spherical Bessel fuuction of the first kiud. while is the dimensionless power per wavenumber octave and P(k.2) is the power spectrum of the mass density fluctuations.,"where is the Fourier transform of the window function $W_R(r)$, and $j_1$ is a spherical Bessel function of the first kind, while is the dimensionless power per wavenumber octave and $P(k,z)$ is the power spectrum of the mass density fluctuations."
 The mean square fluctuation that we would neasure from the actual density contrast depends upon the actual nonlinear power spectrum. but we cau analogously define the linear variance by replacing the power spectrin in the above expression with the linear spectrum. constrained to have the same amplitude for k»0 (i.e. on large scales where nonlinear evolution is neclieible).," The mean square fluctuation that we would measure from the actual density contrast depends upon the actual nonlinear power spectrum, but we can analogously define the linear variance by replacing the power spectrum in the above expression with the linear spectrum, constrained to have the same amplitude for $k\to0$ (i.e., on large scales where nonlinear evolution is negligible)."
 The spatial represcutation of the above integral is Given by the expression where © is the two-point correlation function. aud V—IzR/3.," The spatial representation of the above integral is given by the expression where $\xi$ is the two-point correlation function, and $V = 4\pi
R^3/3$."
 Iu ? we have used the above equation to estimate the true presecut-day value of the oy parameter from the PSC: survey correlation function (?).., In \citet{pairwise} we have used the above equation to estimate the true present-day value of the $\sigma_8$ parameter from the $z$ survey correlation function \citep{hamteg02}.
 We have also used other empirical correlation functions. derived from different survevs. as a template and we found that the resulting value of ay8 was uot sensitive to such variatious: see 7.," We have also used other empirical correlation functions, derived from different surveys, as a template and we found that the resulting value of 8 was not sensitive to such variations; see \citet{pairwise}."
 We use two different methods to estima5 the nonlinear correctious for σε. one based iu perturbation theory. which allows us to express the correction as a simple analytical expressio- and one using a phenomenological mapping based on conservation of pair counts and calibrated musing nunierical smnulatious. which allows us to explore the effects of changing many paranieters iucividuallv.," We use two different methods to estimate the nonlinear corrections for $\sigma_8$, one based in perturbation theory, which allows us to express the correction as a simple analytical expression, and one using a phenomenological mapping based on conservation of pair counts and calibrated using numerical simulations, which allows us to explore the effects of changing many parameters individually."
 Under linear evolution. the spatial aud temporal depeudeuce of clustering separates;," Under linear evolution, the spatial and temporal dependence of clustering separates."
 The density perturbation ó(x.«)=àp/p can be described as (2) where 6) eives the linear density perturbation field as a function of comoving spatial coordinates x at some fiducial tine aud D is the growth function. here parameterized by the scale factor « asa time coordinate (here aud below we keep ouly the fastest-erowing modes).," The density perturbation $\delta({\bf x}, a)=\delta\rho/\rho$ can be described as \citep{pjep}
 where $\delta^{(1)}$ gives the linear density perturbation field as a function of comoving spatial coordinates ${\bf x}$ at some fiducial time and $D$ is the growth function, here parameterized by the scale factor $a$ as a time coordinate (here and below we keep only the fastest-growing modes)."
 In flat ACDAL models the erowing mode is eiven by (?7) where Iu the carly Universe. when the scale factor is sanall. ο20. equation (11)) is well approximated by the expression asin Eiusteim-de Sitter Universe.," In flat $\Lambda$ CDM models the growing mode is given by \citep{Heath:1977p777, pjep2}
 where In the early Universe, when the scale factor is small, $a \to 0$, equation \ref{eqn:Dexact}) ) is well approximated by the expression as in Einstein-de Sitter Universe."
" Iu the opposite ιτ, the cosinologicalOo coustaut becomes dynamically ≼↧∪⋯↕∐⋜⋯↑⋜⋯≼↧↑↕∐∖∐∐↸∖⋜∐⋅∶↴∙⊾↥⋅∪↖↖⇁↑∐↕≯⋜↧↸⊳↑∪↥⋅↖↖⇁↕∐↴∖↴⋜↕⊓∐⋅⋜↧↑↸∖ ⋜↧↑⋜↧↕⊔⋜⋯↕∐∐∐⊔↖⇁⋜↧↕⋯∖∙⋜↕↴∖↴∶↴∙⊾↥⋅⋜↧↖⇁↕↑⋜↧↑↕∪∐⋜↧↕↸⊳↕∏↴∖↴↑↸∖↥⋅↕∐∶↴∙⊾ is balanced bv the effective. force. of accelerated expansion."," In the opposite limit, the cosmological constant becomes dynamically dominant and the linear growth factor will saturate ata maximum value, as gravitational clustering is balanced by the effective force of accelerated expansion."
 It is easv to show that in the liuüt uo. the erowth factor is eiven bw the OXpression ↕↸↴∖↴↕∐∶↴∙⊾↑∐↸∖⋜∏⋝∪↖↽↸∖↑↖↖↽∪⋜↧↴∖↴⋅↖⇁∐∏≻↑∪↑↕↸⊳↸∖↘↻↥⋅↸∖↴∖∷∖↴↕∪∐↴∖↴↖↖⇁↸∖ ∐⋜↧↖↽↸∖↕≯∪∏∐≼↧⋜↧∐↸∖↖↖↽∐⇈↕∐∶↴⋁↕⋟∪↥⋅∐⋯↕⋜↧↕⋟∪↥⋅↑∐↸∖∶↴⋁↥⋅∪↖↖↽↑∐ factor. valid for ⋜↧∐∏⋜↧↑↶∖≼⊲↕≻⋀∖↕↸⊳∪↴∖↴↕⊔∪↕∪∶↴⋁↕↸⊳⋜↧↕ models: In Figue 1 πο show that equation (15)) remains within two-perceut level of the exact solution (11)) iu both the past and the future.," It is easy to show that in the limit $a\to\infty$, the growth factor is given by the expression Using the above two asymptotic expressions we have found a new fitting formula for the growth factor, valid for all flat $\Lambda$ CDM cosmological models: In Figure \ref{fig:Dofa} we show that equation \ref{eqn:Dapprox}) ) remains within two-percent level of the exact solution \ref{eqn:Dexact}) ) in both the past and the future."
 Note that some expressions for D(a) and its logavithinic derivative. dlog D/dloga. frequently quoted in the literature (27). apply ouly to the past and fail for ο> 1.," Note that some expressions for $D(a)$ and its logarithmic derivative, $d \log D/d\log a$ , frequently quoted in the literature \citep{Lahav:1991p833, Carroll:1992p774} apply only to the past and fail for $a>1$ ."
where summation over repeated indices is assumed.,where summation over repeated indices is assumed.
 This Censor must obev a set of svinmetries in order to avoid infinite torques., This tensor must obey a set of symmetries in order to avoid infinite torques.
 In addition we expect it to be symmetric wilh respect of rotations aud reflections around the 2 axis., In addition we expect it to be symmetric with respect of rotations and reflections around the $\hat{z}$ axis.
" With all these svinmetries Ajj, can be shown to have only six independent components.", With all these symmetries $K_{ijmn}$ can be shown to have only six independent components.
 OI those. the particular forms of forcing describedin section 2.2. probe only (wo Bunilies: which will be referred to as A454» and. Aq4344 from now on.," Of those, the particular forms of forcing describedin section \ref{sec: external_forcing} probe only two families: which will be referred to as $K_{1212}$ and $K_{1313}$ from now on."
 To simplify the cliscussion we define uplront the folowing quantities for the z dependent forcing: and for the y dependent forcing: where wr; and y; are the locations of the .c and y collocation grid points. and { are the times at which we have sampled the velocity field.," To simplify the discussion we define upfront the following quantities for the $z$ dependent forcing: and for the $y$ dependent forcing: where $x_i$ and $y_j$ are the locations of the $x$ and $y$ collocation grid points, and $t_w$ are the times at which we have sampled the velocity field."
 We evaluate the sum over an integer number of forcing periods 7., We evaluate the sum over an integer number of forcing periods $T$ .
Infrared divergences limit the perturbative expansion of QCD toO(g’) [9]..,Infrared divergences limit the perturbative expansion of QCD to$O(g^5)$ \cite{Linde}.
 The evaluation up to this order has been carried out [LO] and is found to result in strongly oscillating and hence non-convergent behavior [or 7'€10T;, The evaluation up to this order has been carried out \cite{Arnold} and is found to result in strongly oscillating and hence non-convergent behavior for $T \leq 10~T_c$.
" This has led to considerable efforts to ""repair the difficulty. either by introducing scale effects to allow a svstematic extension of perturbation theory. bevond O(q?) [11]... or by regrouping sets of Feynman diagrams to expand around a ground state including screening effects. (resununed™ perturbation theory |12].. hard. thermal loop approach [13]))."," This has led to considerable efforts to “repair” the difficulty, either by introducing non-perturbative scale effects to allow a systematic extension of perturbation theory beyond $O(g^5)$ \cite{Laine}, or by regrouping sets of Feynman diagrams to expand around a ground state including screening effects (“resummed” perturbation theory \cite{resum}, hard thermal loop approach \cite{HTL}) )."
 In both cases. however. such weak-coupling methods cannot account for the behavior observed in lattice studies.," In both cases, however, such weak-coupling methods cannot account for the behavior observed in lattice studies."
 This holds in particular for $9U(3) gauge theory. where one has results for the continuum limit |1H]:: see TTL (elt) for a comparision to HTL results.," This holds in particular for $SU(3)$ gauge theory, where one has results for the continuum limit \cite{Boyd}; see \\ref{HTL} (left) for a comparision to HTL results."
 It is obvious that no weak-coupling approach can account for the critical behavior near T; (the dip of A(T) as T—11): but also the behavior in the region up to about 5 Τεν with 7?7.N(T)econst.. is not reproduced bv the weak logarithmic form of perturbative studies.," It is obvious that no weak-coupling approach can account for the critical behavior near $T_c$ (the dip of $\Delta(T)$ as $T\to T_c$ ); but also the behavior in the region up to about 5 $T_c$, with $T^2 \Delta(T) \simeq {\rm const.}$, is not reproduced by the weak logarithmic form of perturbative studies."
 The simplest non-perturbative approach. using the bag model form discussed above. does not [are any better.," The simplest non-perturbative approach, using the bag model form discussed above, does not fare any better."
 The bag pressure can be related to the gluon condensate G7 at T=0 (15].. and using numerical estimates for the latter [16].. one obtains a NT) vanishing as T5 and thus much too fast: moreover. the critical dip near T; is also here not given (see (right)).," The bag pressure can be related to the gluon condensate $G^2_0$ at $T=0$ \cite{Leutwyler}, and using numerical estimates for the latter \cite{SVZ}, , one obtains a $\Delta(T)$ vanishing as $T^{-4}$ and thus much too fast; moreover, the critical dip near $T_c$ is also here not given (see \\ref{HTL} (right))."
Flat Spectrum. Racio Quasars (ESI). are. radio. loud Active Galactic ισα (AGN) with broad emission. lines and a non-thermal spectrum. extending from radio to gamma rav energies.,Flat Spectrum Radio Quasars (FSRQ) are radio loud Active Galactic Nuclei (AGN) with broad emission lines and a non-thermal spectrum extending from radio to gamma ray energies.
 They are characterized by luminous core. rapidly. variable non-thermal emission. high radio and optical polarization. Uat-spectrum radio emission. and/or superluminal motion.," They are characterized by luminous core, rapidly variable non-thermal emission, high radio and optical polarization, flat-spectrum radio emission and/or superluminal motion."
 Similar properties are also observed in BL Lacs and these two types of AGN are. classified as blazars., Similar properties are also observed in BL Lacs and these two types of AGN are classified as blazars.
 According to AGN unification scheme. blazars are the class of AGN with a relativistic jet. pointed. close to the line of sight of the observer (Urry&Padovani (1995))).," According to AGN unification scheme, blazars are the class of AGN with a relativistic jet pointed close to the line of sight of the observer \cite{urry95}) )."
 The spectral energy. distribution (SED) of FSRQ is characterized by two broad humps., The spectral energy distribution (SED) of FSRQ is characterized by two broad humps.
 Phe low energy hump in the SED is well understood as svnchrotron emission from a relativistic distribution of electrons., The low energy hump in the SED is well understood as synchrotron emission from a relativistic distribution of electrons.
 Whereas. the origin of high energy. hump is still matter of debate.," Whereas, the origin of high energy hump is still matter of debate."
 There are several models to explain the high energy. emission. based on either leptonic or hadronic interactions (DBlandford.& (1998))).," There are several models to explain the high energy emission based on either leptonic or hadronic interactions \cite{BlandfordLevinson95, 
BloomMarscher96, Aharonian2000, PohlSchlickeiser2000, mannheim98}) )."
 Under the assumption of leptonic models. high energy emission is explained via inverse C'ompton scattering of soft. target. photons.," Under the assumption of leptonic models, high energy emission is explained via inverse Compton scattering of soft target photons."
 The target. photons can be the svnchrotron. photons. (svnehrotron self Compton (SSC)) (Ixonigl.(1981):Alarscher&Cear(1985):GhiselliniAlarasehi (1989))) or the photons external to the jet (external Compton ))(Begelman&Sikora(1987):Melia (1992))).," The target photons can be the synchrotron photons (synchrotron self Compton (SSC)) \cite{konigl81,marscher85,ghisellini89}) ) or the photons external to the jet (external Compton \cite{begelman87,melia89,dermer92}) )."
 These external photons can be the| accretion disk photons (Dermer&Schlickeiser(1993):Boetteher.Alause.&Sehlickeiser (1997))) or the accretion disk photons reprocessed by broad line region. (DLIU) clouds (Sikora. (1996))) or the infrared. radiation. from the dusty torus," These external photons can be the accretion disk photons \cite{dermer93,boettcher97}) ) or the accretion disk photons reprocessed by broad line region (BLR) clouds \cite{sikora94,ghisellini96}) ) or the infrared radiation from the dusty torus"
 These external photons can be the| accretion disk photons (Dermer&Schlickeiser(1993):Boetteher.Alause.&Sehlickeiser (1997))) or the accretion disk photons reprocessed by broad line region. (DLIU) clouds (Sikora. (1996))) or the infrared. radiation. from the dusty torus.," These external photons can be the accretion disk photons \cite{dermer93,boettcher97}) ) or the accretion disk photons reprocessed by broad line region (BLR) clouds \cite{sikora94,ghisellini96}) ) or the infrared radiation from the dusty torus"
simulation to compute total energy. emitted by each Σων particle”.,simulation to compute total energy emitted by each ``N-Body particle”.
" The mass Mg, is given by here fi is the fraction of barvons. 5 is the contribution of barvons to the density parameter and. f,, is the neutral fraction."," The mass $M_{H_I}$ is given by here $f_b$ is the fraction of baryons, $\Omega_b$ is the contribution of baryons to the density parameter and $f_n$ is the neutral fraction."
" We have chosen £3,=0.06 as this value compares well with the observed. abundance of light elements. and primordial nucleosynthesis (Copi.SchrammancTurner 1995).", We have chosen $\Omega_b=0.06$ as this value compares well with the observed abundance of light elements and primordial nucleosynthesis \cite{copnuc}.
. Using this. we can estimate the Ilux received by an observer from an “N-Bocly particle”.," Using this, we can estimate the flux received by an observer from an “N-Body particle”."
" For Q,,=0,1 models the flux contributed. by one particle at. redshift >=3.34 is given by The frequeney width used here corresponds to a velocity dispersion of 200km/s. (Ehe corresponding number for the nmiocdel with 04=0.4 is 0.177,51y.)", For $\Omega_{nr}=\Omega_0 =1$ models the flux contributed by one particle at redshift $z=3.34$ is given by The frequency width used here corresponds to a velocity dispersion of $200$ $/$ s. (The corresponding number for the model with $\Omega_\Lambda = 0.4$ is $0.77\mu Jy$ .)
 In this section we shall outline the results of analysis of the radio maps generated from. N-Bocds simulations., In this section we shall outline the results of analysis of the radio maps generated from N-Body simulations.
 We begin with a pictorial preview of the radio maps., We begin with a pictorial preview of the radio maps.
 Figure 3 shows a sample radio map for each of the three models at redshift =3.34., Figure 3 shows a sample radio map for each of the three models at redshift $z=3.34$.
 The panels of this figure show one frequency channel (chosen to be 125klHlz: which - in velocity units - corresponds to alout 115km /s.) The contours in these racio maps Corresponc to 15. 30 and 605v.," The panels of this figure show one frequency channel (chosen to be $125$ kHz; which - in velocity units - corresponds to about $115$ $/$ s.) The contours in these radio maps correspond to $15$, $30$ and $60\mu$ Jy."
 The pixel size is 3.2 arc minutes zuxl it corresponds to à comoving scale of 24h ‘Alpe for mocdels E and LE., The pixel size is $3.2$ arc minutes and it corresponds to a comoving scale of $2.7h^{-1}$ Mpc for models I and II.
 (3.55 “Alpe for moclel LLL), $3.5h^{-1}$ Mpc for model III.)
 1t is clear from tjese panels that models E (sCDAL) and. 111 (LCD) have comparable signal whereas model LE (NDM). as it has less power at smaller scales. has somewhat lower," It is clear from these panels that models I (sCDM) and III (LCDM) have comparable signal whereas model II (MDM), as it has less power at smaller scales, has somewhat lower"
1991: Waleffe1997: Hamiltonefaf1995: see section ??)).,; \citealt{Wal97}; ; \citealt{Ham95}; see section \ref{plcou}) ).
 The fact that [ree rotating lavers aud Couette-Tavlor flows remain turbulent at larger levels of rotation than the ones to which turbulence is lost in these simulations implies (hat a different mechanism for sustaining (turbulence is at work in these flows. and that it operates at scales comparable to. but apparently smaller than. the estimate of Eq. (17)).," The fact that free rotating layers and Couette-Taylor flows remain turbulent at larger levels of rotation than the ones to which turbulence is lost in these simulations implies that a different mechanism for sustaining turbulence is at work in these flows, and that it operates at scales comparable to, but apparently smaller than, the estimate of Eq. \ref{lm}) )."
 This other mechanism has not vet been found in munerical simulations., This other mechanism has not yet been found in numerical simulations.
 It would be interesting to know whether this change of mechanism is related to the fact that the Coriolis force apparently selects the direction of instability of finite amplitude defects (Johnson1963)., It would be interesting to know whether this change of mechanism is related to the fact that the Coriolis force apparently selects the direction of instability of finite amplitude defects \citep{John63}.
. The same line of argument. applies to the shearing sheet simulations of Balbus and Hawlevefaf.(1999)., The same line of argument applies to the shearing sheet simulations of \citet{BHS96} and \citet{HBW99}.
. Indeed. the effective Revnolds number of these simulations is not an issue. as the code used bv Lawleyefa£(1995) and Lawleyefa£.(1999) is able to find turbulence — or at least the large scale mechanism already alluded to in non-rotating Couette flows. and this happens only for Revnolds numbers larger than at least 1500.," Indeed, the effective Reynolds number of these simulations is not an issue, as the code used by \citet{HGB95} and \citet{HBW99} is able to find turbulence — or at least the large scale mechanism already alluded to — in non-rotating Couette flows, and this happens only for Reynolds numbers larger than at least $1500$."
 Also. (he argument developed in section ?? shows that the Coriolis force bv itself should not change ihe minimal Revuolds number for the onset of turbulence. an inference confirmed by the [act that turbulence is seen developing in the rotating [ος shear laver of experiments of Bicokhti [or roughly comparable Reynolds numbers.," Also, the argument developed in section \ref{turbord} shows that the Coriolis force by itself should not change the minimal Reynolds number for the onset of turbulence, an inference confirmed by the fact that turbulence is seen developing in the rotating free shear layer of experiments of \citet{Bid92} for roughly comparable Reynolds numbers."
 Under the assumption (cf (he arguments developed above) that Eq. (17)), Under the assumption (cf the arguments developed above) that Eq. \ref{lm}) )
" provides an estimate for (he largest turbulent scale which is overestimated by a [actor of at least 3 in the presence of a Coriolis force term. one obtains fy,<Ay/100. with (possibly much) smaller values more than likely."," provides an estimate for the largest turbulent scale which is overestimated by a factor of at least $3$ in the presence of a Coriolis force term, one obtains $l_M\lesssim
\Delta y/100$, with (possibly much) smaller values more than likely."
 This is most probably (oo close to the largest resolution achieved in the shearing sheet simulations (ihe number of zones in anv clirection being no larger than 250). especially when artificial viscosity is taken into account. for turbulence to show up in these simulations.," This is most probably too close to the largest resolution achieved in the shearing sheet simulations (the number of zones in any direction being no larger than 250), especially when artificial viscosity is taken into account, for turbulence to show up in these simulations."
 To conclude this section. it is worth noting that some other numerical questions must be considered to find turbulence in these systems.," To conclude this section, it is worth noting that some other numerical questions must be considered to find turbulence in these systems."
 First. it is well known from experiments wilh subcritical flows that the way perturbations of the flow are designed has an influence on (he appearance of turbulence.," First, it is well known from experiments with subcritical flows that the way perturbations of the flow are designed has an influence on the appearance of turbulence."
 This suggests that some care must be exercised in the choice of the initial conditions in numerical experiments; in particular. it might be uselul io ensure that at least some condition of finite amplitude instability is satisfied in this choice.," This suggests that some care must be exercised in the choice of the initial conditions in numerical experiments; in particular, it might be useful to ensure that at least some condition of finite amplitude instability is satisfied in this choice."
 Secondly. the role of the choice of the Courant number is not completely obvious. even in situations where the CFL condition is not violated.," Secondly, the role of the choice of the Courant number is not completely obvious, even in situations where the CFL condition is not violated."
 For example. in a series of vel unpublished simulations of linearly stable Couette-Tavlor flows performed with the Zeus code in collaboration with David Clarke. we did initially find that turbulence would set in [or flows which are “not too far” [rom planar Couette flows (including some roughly. keplerian flows). but it would eventually disappear when reducing the maximal allowed. time-step. allhough the CFL condition was satisfied in all runs.," For example, in a series of yet unpublished simulations of linearly stable Couette-Taylor flows performed with the Zeus code in collaboration with David Clarke, we did initially find that turbulence would set in for flows which are “not too far"" from planar Couette flows (including some roughly keplerian flows), but it would eventually disappear when reducing the maximal allowed time-step, although the CFL condition was satisfied in all runs."
 The reason of this behavior is not vel completely elucidated. but it appearsto have some direct connection to (he question," The reason of this behavior is not yet completely elucidated, but it appearsto have some direct connection to the question"
There are two methods of testing the source of illumination of Barnard's Loop.,There are two methods of testing the source of illumination of Barnard's Loop.
 The first test is to determine if the total stellar luminosity in the LyC that is derived from the surface brightness is consistent with the available flux from the enclosed stars and is considered in 33.2., The first test is to determine if the total stellar luminosity in the LyC that is derived from the surface brightness is consistent with the available flux from the enclosed stars and is considered in 3.2.
 The second test is to determine if the observed emission-ine ratios can be explained by photoionization caused by the enclosed stars and is considered in 33.4 for the well studied object M 43 as a test object and then in 33.5 we apply this method to the Barnard's Loop., The second test is to determine if the observed emission-ine ratios can be explained by photoionization caused by the enclosed stars and is considered in 3.4 for the well studied object M 43 as a test object and then in 3.5 we apply this method to the Barnard's Loop.
" The observations of the emission distribution, the surface brightness, and the known distance allow us to determine some basic properties of Barnard's Loop."," The observations of the emission distribution, the surface brightness, and the known distance allow us to determine some basic properties of Barnard's Loop."
" Fortunately the interstellar extinction in the region is known to be negligible (Peimbertetal. 2006),, a conclusion consistent with the fact that our new flux ratio of 2.90 is that predicted for an unreddened gas of about 9000 K. In their fundamental study of the Orion Nebula, Baldwinetal.(1991) demonstrated that the surface brightness in the rrecombination line along a line of sight from the observer to the photoionizing star would be proportional to the flux of ionizing photons ¢(H) in units of photons cm? $1 reaching a thin slab of gas."," Fortunately the interstellar extinction in the region is known to be negligible \citep{pei75,mad06}, a conclusion consistent with the fact that our new flux ratio of 2.90 is that predicted for an unreddened gas of about 9000 K. In their fundamental study of the Orion Nebula, \citet{b91} demonstrated that the surface brightness in the recombination line along a line of sight from the observer to the photoionizing star would be proportional to the flux of ionizing photons $\phi$ (H) in units of photons $^{-2}$ $^{-1}$ reaching a thin slab of gas."
 This approximation is probably valid for Barnard's Loop because the low inner-region densities would not be expected to absorb LyC photons., This approximation is probably valid for Barnard's Loop because the low inner-region densities would not be expected to absorb LyC photons.
" If one views the emitting layer face-on, then the relation would be 9(H)—47(ap/a115)S(H8)) where the surface brightness S(H)) is expressed in aand ap and ans are respectively the total recombination coefficient for hydrogen and the effective recombination coefficient for the"," If one views the emitting layer face-on, then the relation would be $\phi$ $\pi$ $\rm \alpha _{B}$ $\rm \alpha ^{eff}_{H\beta}$ ) where the surface brightness ) is expressed in and $\rm \alpha _{B}$ and $\alpha ^{eff}_{H\beta}$ are respectively the total recombination coefficient for hydrogen and the effective recombination coefficient for the"
lt is convenient to subtract the selection function contribution: We now construct the combined density field (1)=(1f)nCr)SOF) and write (for the purposes of. this calculation) a simple power spectrum estimator where md—í(linci=7ueUk).,"It is convenient to subtract the selection function contribution: We now construct the combined density field $m(\vec{x}) = (1-f)
n(\vec{x}) + S(\vec{x})$ and write (for the purposes of this calculation) a simple power spectrum estimator where $\tilde{m}'(\vec{k}) = (1-f) \tilde{n}'(\vec{k}) +
\tilde{S}'(\vec{k})$."
" The terms needed to determine. |m'(A)]|7u ere: where Writing the convolved power spectrum as 4(&)= CU. the final result. for the power spectrum. estimate is: In the special case that the selection. function. and scattering functions are constant we find that: IL the Fourier modes. & are binned in spherical shells as à function of &=IK then. as the value of & increases. the contribution of the mode with &,=0 becomes less important and {λα}=(1[YPU) is an increasingly &ood approximation."," The terms needed to determine $|\tilde{m}'(k)|^2$ are: where Writing the convolved power spectrum as $P_c(\vec{k}) =
|\tilde{C}(\vec{k})|^2$ , the final result for the power spectrum estimate is: In the special case that the selection function and scattering functions are constant we find that: If the Fourier modes $\vec{k}$ are binned in spherical shells as a function of $k = |\vec{k}|$ then, as the value of $k$ increases, the contribution of the mode with $k_x = 0$ becomes less important and $P_{\rm est}(k) = (1-f)^2 P(k)$ is an increasingly good approximation."
 We now consider the case where a fraction f of redshifts are scattered by a fixed. displacement along the w-axis. aa convolution of the density field by a funetion C.," We now consider the case where a fraction $f$ of redshifts are scattered by a fixed displacement along the $x$ -axis, a convolution of the density field by a function $U$."
" Equation 26 is replaced by the relation with the Fourier transform Defining SR)=SUA)FNU(,TT we find that For à simple shift Cr)οro) we find that These formulae are only approximations in the case of real data."," Equation \ref{eqsx} is replaced by the relation with the Fourier transform Defining $\tilde{S}'(\vec{k}) = \tilde{S}(\vec{k}) - f N \tilde{U}(k_x)
\tilde{W}(\vec{k})$ we find that For a simple shift $U(x) = \delta(x-x_0)$ we find that These formulae are only approximations in the case of real data."
 Firstly. the sky is curved and the redshift scatters do not happen along a single axis of the cube.," Firstly, the sky is curved and the redshift scatters do not happen along a single axis of the cube."
 Secondly. the blunder fraction depends on redshift.," Secondly, the blunder fraction depends on redshift."
 Phirdlv. the blunders due to line confusion are not a strict convolution of the density field. but a transformation in redshift τι»2o of the form so=C'(1|zi)1l. where €' is a constant depending on the rest wavelengths of the lines.," Thirdly, the blunders due to line confusion are not a strict convolution of the density field, but a transformation in redshift $z_1 \rightarrow z_2$ of the form $z_2 = C (1+z_1) - 1$, where $C$ is a constant depending on the rest wavelengths of the lines."
" Fourthly. if à convolved redshift’ is shifted: bevond the edge of the density cube. it does not “wrap around"" as required by periodic boundary conditions."," Fourthly, if a convolved redshift is shifted beyond the edge of the density cube, it does not “wrap around” as required by periodic boundary conditions."
 Fifthlv. FINP estimation of the power spectrum is used rather than Equation 32..," Fifthly, FKP estimation of the power spectrum is used rather than Equation \ref{eqpkestsimp}."
 In order to measure the distortion of our measured power spectrum due to redshift blunders for the real data. we therefore. created Monte Carlo simulations of galaxy catalogues with a known input power spectrum and the same selection function. WCF) as each of our survey regions.," In order to measure the distortion of our measured power spectrum due to redshift blunders for the real data, we therefore created Monte Carlo simulations of galaxy catalogues with a known input power spectrum and the same selection function $W(\vec{x})$ as each of our survey regions."
 We then applied the redshift blunder distributions of Figures 1l and 12 to the mock catalogues ancl re-measured the power spectrum., We then applied the redshift blunder distributions of Figures \ref{figzblundline} and \ref{figzblundsky} to the mock catalogues and re-measured the power spectrum.
 Comparison of the input and output power spectra. averaging over many Monte. Carlo. realizations. provided. the correction. factor due to redshift bluncders.," Comparison of the input and output power spectra, averaging over many Monte Carlo realizations, provided the correction factor due to redshift blunders."
 We tested our code by reproducing the relations given in Equations 37. and 42. in the [lat-sky. case., We tested our code by reproducing the relations given in Equations \ref{eqpkblund1} and \ref{eqpkblund2} in the flat-sky case.
" ligure 13. plots the correction factors for the ""angle-averaged"" power spectrum. P(&) Lor cach of the survey regions (taking a redshift interval 0.3.<2. 0.9).", Figure \ref{figpkbadzcorr} plots the correction factors for the ``angle-averaged'' power spectrum $P(k)$ for each of the survey regions (taking a redshift interval $0.3 < z < 0.9$ ).
 We see that a good approximation for the correction for scales kc0.05h Alpe 4.is à constant. although a more significant⋠⋠⋅ correction is required for large scales &«0.057 +.," We see that a good approximation for the correction for scales $k > 0.05 \, h$ $^{-1}$ is a constant, although a more significant correction is required for large scales $k < 0.05 \, h$ $^{-1}$."
 For koc0.055 the approximately constant. correction [actor is not exactly equal to (1f) as predicted. by Equation 39.. where f is the average redshift’ blunder rate of the catalogue: this is due to the mis-estimation of the denominator of Equation 13. that occurs because VCr) is determined from the galaxy redshift distribution including blunders. as cleseribecl in Section. 2.5..," For $k > 0.05
\, h$ $^{-1}$ the approximately constant correction factor is not exactly equal to $(1-f)^{-2}$ as predicted by Equation \ref{eqpkblundsimp}, where $f$ is the average redshift blunder rate of the catalogue; this is due to the mis-estimation of the denominator of Equation \ref{eqpkest} that occurs because $W(\vec{x})$ is determined from the galaxy redshift distribution including blunders, as described in Section \ref{secnz}."
 The Monte. Carlo simulations performed. here also correct. for this small bias in power spectrum estimation., The Monte Carlo simulations performed here also correct for this small bias in power spectrum estimation.
 In this study we analyzed a galaxy sample drawn. from WieeleZ survey observations prior to July 2009 in SDSS regions of our optical imaging (9-hr. Li-hr. 15-hr).," In this study we analyzed a galaxy sample drawn from WiggleZ survey observations prior to July 2009 in SDSS regions of our optical imaging (9-hr, 11-hr, 15-hr)."
 Figure 14 plots the (RA... Dee.) distribution of these redshifts.," Figure \ref{figradec} plots the (R.A., Dec.) distribution of these redshifts."
 We, We
a normal error distribution vields Siuce the true parallax is hekl fixed at ap assumed value. the proper notion aud magnitude constraints do not affect this factor.,"a normal error distribution yields Since the true parallax is held fixed at an assumed value, the proper motion and magnitude constraints do not affect this factor."
 The secoud [actor is thepriori probajlity of a particuar parallax. given ouly the prior iuformation.," The second factor is the probability of a particular parallax, given only the prior information."
 This itself is composed of sever:il factexs., This itself is composed of several factors.
 The proper motion probability deusity is included here., The proper motion probability density is included here.
 TIe inaguitucde constraint is solrewha more ciffieult. because of bias.," The magnitude constraint is somewhat more difficult, because of bias."
 treats the case of a Caussian distri»utio1 of absolute iuagnitudes for the type in question. aud fornulates a Maluquist-type adjustineut to tle WOs likely absolute magnitudee which accounts for the tendency to dick out. absolutely uxigher (heice more distant) members of a population with €a non-zero luitosity dispersion.," \citet{smith87b} treats the case of a Gaussian distribution of absolute magnitudes for the type in question, and formulates a Malmquist-type adjustment to the most likely absolute magnitude which accounts for the tendency to pick out absolutely brighter (hence more distant) members of a population with a non-zero luminosity dispersion."
 The cor'ection rejxaces the mean absolute magnitude Aly with AL*=Aly--L&loy)., The correction replaces the mean absolute magnitude $M_0$ with $M^* = M_0 - 1.84 \sigma_M^2$.
 Alternatively. οle Ca1 formulae a deusity by simply multiplying tle volume elenent (xl/z ) by the appropriae Gaussian welelitineg [uuctiou ceutered on Adj.," Alternatively, one can formulate a density by simply multiplying the volume element $\propto 1/\pi^4$ ) by the appropriate Gaussian weighting function centered on $M_0$."
 Somewhat COLnter-intuitivey. these approaches give the same orobability deusity.," Somewhat counter-intuitively, these approaches give the same probability density."
 Because of the very broad Ciatussialis used ο Characterize lie uuluosltv p‘iors of most of the cataclysimics. these fuuctions eid up resemibliο pure 1/24 distribloli. excert that the singularity as z—0 is eliminated by je Caussial cuolf in absolute hagitucle.," Because of the very broad Gaussians used to characterize the luminosity priors of most of the cataclysmics, these functions end up resembling pure $1/\pi^4$ distributions, except that the singularity as $\pi \rightarrow 0$ is eliminated by the Gaussian cutoff in absolute magnitude."
 The calcu€aion for each star p'oceeded as follows., The calculation for each star proceeded as follows.
" The estimated a,j, aud its external error. je estmatec :ibsolute magnituc ead catalogued apparent maguitude (as appropriate for the type ) [CV ancl outrst state) we'e ablated: the values used are egiven in Table f£ aud commented oi urtler in the uees )elow."," The estimated $\pi_{\rm abs}$ and its external error, the estimated absolute magnitude and catalogued apparent magnitude (as appropriate for the type of CV and outburst state) were tabulated; the values used are given in Table 4 and commented on further in the notes below."
 A erid of drte parallax values z was constructed. from 0.1 to 30 mas jiu Q.1 mas inc'e11011s.," A grid of `true' parallax values $\pi$ was constructed, from 0.1 to 30 mas in 0.1 mas increments."
 This ipper init was Chosen to be safely larger than any of the measure IATAlaxes., This upper limit was chosen to be safely larger than any of the measured parallaxes.
 At οιch j»arallax. the p‘obaility density {ΟΗΕ} was computed. and a cumulative [1istribution fux‘tion was forned from these.," At each parallax, the probability density $P(D|HI)$ was computed, and a cumulative distribution function was formed from these."
 The poluts at which the cumulative cdistributior equaed 0.50. 0.159. aud 0.8LL were take Las the best estimate of the parallax and the positive aux 1egaive I-lgma! error bars.," The points at which the cumulative distribution equaled 0.50, 0.159, and 0.841 were taken as the best estimate of the parallax and the positive and negative `1-sigma' error bars."
 Fig., Fig.
 1. illtstrates this process for VY Aqr. aud the last columns of Table 3 give the results.," 1 illustrates this process for VY Aqr, and the last columns of Table 3 give the results."
 Table [ πιαίσος the parallax. measu‘ements aid the distances derived. [rom them., Table 4 summarizes the parallax measurements and the distances derived from them.
 A discussion of individual objects follows., A discussion of individual objects follows.
" The parallax alone. zi,=11.2drl. Linas. gives a distauce uear 89 pc."," The parallax alone, $\pi_{\rm abs} = 11.2 \pm 1.4$ mas, gives a distance near 89 pc."
 The relative error is small enough that the Bayesian adjustineuts to this are fairly minor., The relative error is small enough that the Bayesian adjustments to this are fairly minor.
 VY Aqr is an SU UMa star with au orbital period of 0.06309(1) d ((Thorstenseun&Taylor1997)., VY Aqr is an SU UMa star with an orbital period of 0.06309(4) d \citep{uvvyv1504}.
. The o'bital inclination is uuknowu. but emission lines in quiescence are strouely double-peaked. suggestin:[n]) i>20 degrees. aud there is no lint of au eclipse. suggesting 7«75 degrees. so I adopt i—6341κ)+) degrees. which," The orbital inclination is unknown, but emission lines in quiescence are strongly double-peaked, suggesting $i > 50$ degrees, and there is no hint of an eclipse, suggesting $i < 75$ degrees, so I adopt $i = 63 \pm 13$ degrees, which"
this context.,this context.
 Even for such a simple svstem. we then ask the In a companion paper (191 Hajj. Monneau |?]}). we will provide some partial answers to (his question.," Even for such a simple system, we then ask the In a companion paper (El Hajj, Monneau \cite{EM2}) ), we will provide some partial answers to this question."
 Let us first consider a smooth function 4=(ul...u). solution of the following non-conservative hyperbolic system: where the space of states U is now an open subset of EZ. and for each uw. F(a) is a (dxd)- and the map F is of class C*(U).," Let us first consider a smooth function $u=(u^1,\dots,u^d)$, solution of the following non-conservative hyperbolic system: where the space of states $U$ is now an open subset of $\R^d$ , and for each $u$, $F(u)$ is a $(d\times d)$ -matrix and the map $F$ is of class $C^1(U)$."
 The system (1.11)) is said (dxd) hyperbolic. if EQ) has d real eigenvalues and is diagonalizable lor anv given uv on the domain under consideration.," The system \ref{EM:lef}) ) is said $(d\times d)$ hyperbolic, if $F(u)$ has $d$ real eigenvalues and is diagonalizable for any given $u$ on the domain under consideration."
 By definition. such a system is said to be diagonalizable. if there exists a smooth transformation ie2Gel(u).....w(u)) with non-vanishing Jacobian such that (1.11)) can be equivalently rewritten(lor smooth solutions) as the following svstem where A’ are smooth functions of «e.," By definition, such a system is said to be diagonalizable, if there exists a smooth transformation $w = (w^1(u),\dots,w^d(u))$ with non-vanishing Jacobian such that \ref{EM:lef}) ) can be equivalently rewritten(for smooth solutions) as the following system where $\l^i$ are smooth functions of $w$."
 Such functions iw! are called. strict ;- Riemann Invariant., Such functions $w^i$ are called strict $i$ -Riemann invariant.
 Qur approach can give continuous solutions to the diagonalized svstem. which provided continuous solution to the original svstem (1.11)).," Our approach can give continuous solutions to the diagonalized system, which provided continuous solution to the original system \ref{EM:lef}) )."
 For a scalar conservation law. which corresponds (o svstem (1.11)) in the case d=1 where Fu)=h'(u) is the derivative of some flux function A. the global existence and uniqueness of BV. solutions has been established by Oleinik ΤΕ in one space dimension.," For a scalar conservation law, which corresponds to system \ref{EM:lef}) ) in the case $d=1$ where $F(u)= h'(u)$ is the derivative of some flux function $h$, the global existence and uniqueness of $BV$ solutions has been established by Oleinik \cite{Ole} in one space dimension."
 The famous paper of INruzhkov |?| covers the more general class of L* solutions. in several space dimensions.," The famous paper of Kruzhkov \cite{Kru} covers the more general class of $L^{\infty}$ solutions, in several space dimensions."
 For an alternative approach based on thenotion of entropy process solutions. see for instance Evinarel et al. |?]..," For an alternative approach based on thenotion of entropy process solutions, see for instance Eymard et al. \cite{Eymard}. ."
 For a different approach based on a kinetic formulation. see also Lions et al. [?]..," For a different approach based on a kinetic formulation, see also Lions et al. \cite{LBT}. ."
shows the results for the 3 spheroidal models: 1erneatulst: PlIumnxv: and Lowered Evans.,shows the results for the 3 spheroidal models: Hernquist; Plummer; and Lowered Evans.
 Each model is populated with 20OOO particles anc the energy change is fractional cllferenc “cof the particle energy at times / and fo., Each model is populated with 20000 particles and the energy change is fractional difference of the particle energy at times $t$ and $t_0$.
 One will notice aIl three SCE particle plots exhibit a vertical baud slructure at certain binding energies., One will notice all three SCF particle plots exhibit a vertical band structure at certain binding energies.
 The cause of some particles being more liable to undergo two-bods. relaxation than otier is mace clearer by Figure(6)) which shows the same plots for three Lowered Evans models with 0. 500. and 100 Vtree particles respectively.," The cause of some particles being more liable to undergo two-body relaxation than other is made clearer by \ref{pen2}) ) which shows the same plots for three Lowered Evans models with 0, 500, and 1000 tree particles respectively."
 The greater the fraction of tree yarticles the wider and. more distinct the vertical features that appear in the regions of low binding energies., The greater the fraction of tree particles the wider and more distinct the vertical features that appear in the regions of low binding energies.
 The first point to bear in mind is that at. low energies the fracional energy changes will become laree because the bincling energy is approaching zero., The first point to bear in mind is that at low energies the fractional energy changes will become large because the binding energy is approaching zero.
 Secondly it is clear that greater the num»er of tree particles the greater the fractional CHCLEN change., Secondly it is clear that greater the number of tree particles the greater the fractional energy change.
 This is cue to the fact that individual SCE xwticles interact directly with the tree particles. and πο he more tree xuticles the greater the energy. exchange otween the two components.," This is due to the fact that individual SCF particles interact directly with the tree particles, and so the more tree particles the greater the energy exchange between the two components."
 Finally the bands that are Seen appear as a result of truncating the SCT expansion., Finally the bands that are seen appear as a result of truncating the SCF expansion.
 The expanded potenial of the svstem is derived from the actual OSLions of all the constituent SCE particles., The expanded potential of the system is derived from the actual positions of all the constituent SCF particles.
 The actual ining cnerev of a particle may not agree exactly with he OLEntial given by the expansion. causing an error in he LOsuting acceleration.," The actual binding energy of a particle may not agree exactly with the potential given by the expansion, causing an error in the resulting acceleration."
 There will be certain values of bindine enerον which correspond to the divergences in the truncaed expanded potential from the true value. thus the partic‘los al rose binding energies will experience greater errors in acceleration.," There will be certain values of binding energy which correspond to the divergences in the truncated expanded potential from the true value, thus the particles at those binding energies will experience greater errors in acceleration."
 Clearly. the presence of Tree. particles in he, Clearly the presence of Tree particles in the
For the sake of illustration. assume that the parameter 6; is well constrained by some data.d. so that the posterior probability p(8;|d.72;) is well-peaked within the prior range. where {/{ is a step function that enforces the prior range and we have chosen a fla? prior on 6). 7(4;)=1/2N06;. within this range.,"For the sake of illustration, assume that the parameter $\theta_j$ is well constrained by some data, so that the posterior probability $p(\theta_j|{\bf d},\mathcal{H}_i)$ is well-peaked within the prior range, where $H$ is a step function that enforces the prior range and we have chosen a prior on $\theta_j$, $\pi(\theta_j) = 1/\Delta \theta_j$, within this range."
 Since the posterior distribution is well-peakect. we can approximate the integral Eq. (5))," Since the posterior distribution is well-peaked, we can approximate the integral Eq. \ref{evidence2}) )"
 using Laplace's method with Eq. (4)):, using Laplace's method with Eq. \ref{bayes}) ):
 we simply multiply the height of the un-normatizecl posterior (the numerator in Ίσα. (4))). p(8;|]d.Hi).," we simply multiply the height of the un-normalized posterior (the numerator in Eq. \ref{bayes}) )), $\tilde{p}(\bar{\theta}_{j}|{\bf d},\mathcal{H}_i)$,"
 by its width. 80jhe where 6; denotes the point of maximun likelihood.," by its width, $\sigma_{\theta_j|d}$, where $\bar{\theta}_{j}$ denotes the point of maximum likelihood."
 In terms of the likelihood function this becomes. While only an approximation. this expression nicely reveals the essential ingredients of Bayesian mocel selection: a hieh-valuecd maximum likelihood clearly increases the evidence in favor of the model. while thefactor. ]—ay)XO;l. penalizes overly complex or poorly predictive models.," In terms of the likelihood function this becomes, While only an approximation, this expression nicely reveals the essential ingredients of Bayesian model selection: a high-valued maximum likelihood clearly increases the evidence in favor of the model, while the, $\beta
= {\sigma_{\theta_j|d}}/{\Delta \theta_j} \leq 1$, penalizes overly complex or poorly predictive models."
 Models with a prior volume much larger than the posterior. volume. οκ1. are not considered predictive because they can accommodate a wide range of parameter values before the data is collected.," Models with a prior volume much larger than the posterior volume, $\beta \ll 1$, are not considered predictive because they can accommodate a wide range of parameter values before the data is collected."
 Complex niodels with loose priors and many [ree parameters that are well-constrainecdd by the cata are therefore penalized by the Occam factor. anc will consequently. have lower evidence than a simpler. more predictive model that fits the cata equally well.," Complex models with loose priors and many free parameters that are well-constrained by the data are therefore penalized by the Occam factor, and will consequently have lower evidence than a simpler, more predictive model that fits the data equally well."
 We are now ready to co model selection: we simply compute the Bavesian evidence for cach mocel and compare., We are now ready to do model selection: we simply compute the Bayesian evidence for each model and compare.
 Given two competing models. Ho and δι. this can be done via the Baves factor. A rubric for scoring thesignificance of a model is eiven by the well-known Jelrevs’ scale (Jellrevs1961).," Given two competing models, $\mathcal{H}_0$ and $\mathcal{H}_1$, this can be done via the Bayes factor, A rubric for scoring thesignificance of a model is given by the well-known Jeffreys' scale \citep{Jeffreys}."
. The scale rates: [InBor}<1 (indecisive). 1<[lnBu]2.5 (substantial). 2.5<|InBor)«5 (strong). and |InBor}>5 (decisive). with InBor-:0 (InBor« 0) favoring Ho (H4).," The scale rates: $|\ln B_{01}| < 1$ (indecisive), $1<|\ln B_{01}| < 2.5$ (substantial), $2.5<|\ln B_{01}| < 5$ (strong), and $|\ln B_{01}| > 5$ (decisive), with $\ln B_{01} > 0$ $\ln B_{01} < 0$ ) favoring $\mathcal{H}_0$ $\mathcal{H}_1$ )."
 In this work. we will quote results for strong and. decisive evidence. corresponding to odds ratios of 12:1 and 150:1. respectively.," In this work, we will quote results for strong and decisive evidence, corresponding to odds ratios of 12:1 and 150:1, respectively."
 While the Laplace method. is instructive. it is. not sullicient for an accurate determination. of the evidence.," While the Laplace method is instructive, it is not sufficient for an accurate determination of the evidence."
 While evaluating the integral in Eq. (5)), While evaluating the integral in Eq. \ref{evidence2}) )
 is computationally demancing. various methods have been applied to problems of model selection in astrophysics. including nested sampling (Skilling2004:Mukherjeeetal.2006). and thermodynamic integration (OItuanaidhetal.1996:Beltran 2005).," is computationally demanding, various methods have been applied to problems of model selection in astrophysics, including nested sampling \citep{Skilling,Mukherjee:2005wg} and thermodynamic integration \citep{Obook,Beltran}. ."
. Ilere we make use of theratio (Dickey 1971).. an exact. analvtical expression for Do; that can be applied: whenever the models to be compared. are nested.," Here we make use of the \citep{Dickey}, an exact analytical expression for $B_{01}$ that can be applied whenever the models to be compared are nested."
 Alodels Hy ancl Hy share the same cosmological paranicters. except for mp which is set to zero in Hy.," Models $\mathcal{H}_0$ and $\mathcal{H}_1$ share the same cosmological parameters, ${\bm \psi}$, except for $n_T$ which is set to zero in $\mathcal{H}_0$."
" For separable priors. z(ab|mr.Hinoπα[Hy). the Bayes factor takes the form where η].Ai),-0 is the marginalized posterior probability of n. under model Hy. evaluated at mp=0."," For separable priors, $\pi({\bm \psi}|n_T, \mathcal{H}_1)|_{n_T = 0} = \pi({\bm \psi}|\mathcal{H}_0)$, the Bayes factor takes the form where $p(n_T|{\bf d},\mathcal{H}_1)|_{n_T = 0}$ is the marginalized posterior probability of $n_T$ under model $\mathcal{H}_1$, evaluated at $n_T = 0$."
" This quantity can be obtained relatively easily using Markov Chain Monte Carlo (AICAIC) techniques. and. in principle. gives Lo, as a function. of Dp."," This quantity can be obtained relatively easily using Markov Chain Monte Carlo (MCMC) techniques, and, in principle, gives $B_{01}$ as a function of $n_T$."
 This is the function that we seek to determine in this analwsis. for a variety of current anc proposed. CALB experiments.," This is the function that we seek to determine in this analysis, for a variety of current and proposed CMB experiments."
" For each experiment. we will obtain projections by generating simulated. CMD data across a range of values of ny. and determine how the evidence for FH, builds as [n7| grows."," For each experiment, we will obtain projections by generating simulated CMB data across a range of values of $n_T$, and determine how the evidence for $\mathcal{H}_1$ builds as $|n_T|$ grows."
" Primordial perturbations impart inhomogencities in the photon temperature at decoupling. measured. today as directional anisotropies on the last scattering sphere. where the multipole moments. (;,,. are complex Gaussian random variables with variance (αλαὃν)=οδιὃν in the directionn."," Primordial perturbations impart inhomogeneities in the photon temperature at decoupling, measured today as directional anisotropies on the last scattering sphere, where the multipole moments, $a_{\ell m}$, are complex Gaussian random variables with variance $\langle a^{T*}_{\ell m}a^T_{\ell' m'} \rangle = C^{TT}_\ell \delta_{\ell
\ell'}\delta_{m m'}$ in the direction."
" Quacrupolar temperature anisotropies at decoupling (and again at reionization) are projected into anisotropies in the polarization of the CM and can be similarlv. cecomposed (Ixamionkowskietal. L997). where the Y,nra777 are clectric-tvpe (curl-free) and magnetic-type (divergencee-free) tensor spherical harmonics. respectively."," Quadrupolar temperature anisotropies at decoupling (and again at reionization) are projected into anisotropies in the polarization of the CMB and can be similarly decomposed \citep{Kamionkowski:1996ks}, where the $Y^{E,B}_{(\ell m)ab}$ are electric-type (curl-free) and magnetic-type (divergence-free) tensor spherical harmonics, respectively."
 Ehe polarization anisotropies are described by the correlations. and One nonzero ο correlation. Primordial density perturbations can be constrained by measurements. of the temperature anc E-miocle polarization anisotropies. while primordial gravitationalwaves additionally create a B-moce polarization pattern (IKamionkowskictal.1997:Seljak&Zaldarriaga 1997).," The polarization anisotropies are described by the correlations, and one nonzero cross correlation, Primordial density perturbations can be constrained by measurements of the temperature and E-mode polarization anisotropies, while primordial gravitationalwaves additionally create a B-mode polarization pattern \citep{Kamionkowski:1996zd,Seljak:1996gy}. ."
. The D-mocde signal is therefore a key indicator ofprimordial lensors., The B-mode signal is therefore a key indicator ofprimordial tensors.
 Our projections are based on simulated datasets., Our projections are based on simulated datasets.
" Since we do not have access to the true distribution of the eai,s. an estimator is formed from their measured. values."," Since we do not have access to the true distribution of the $a_{\ell m}$ 's, an estimator is formed from their measured values,"
Levensonetal. 20063).,\citealt{LE06.1}) ).
 The Compton-thick scenario is also consistent with the low coluun deusity of the galaxy. siuce the column density nieasureimoeut presupposed that the source was not Compton thick.," The Compton-thick scenario is also consistent with the low column density of the galaxy, since the column density measurement presupposed that the source was not Compton thick."
 We conclude that SDSS J171511.05|G008S35.7. most likely hosts Comptou-thick dual AGN. because it is the scenario that is most consistent with the existing data.," We conclude that SDSS J171544.05+600835.7 most likely hosts Compton-thick dual AGN, because it is the scenario that is most consistent with the existing data."
 AGN jets melt also explain the observations. and deeper X-ray observations could distinguish between these two possibilities.," AGN jets might also explain the observations, and deeper X-ray observations could distinguish between these two possibilities."
 These more sensitive X-ray mneasurenments would cuable a test of the Compton-thick scenario. siuce better spectral measurements in the soft N-rav baud could provide a possible measurement of the Fe K line. which is secu in heavily absorbed AGN (for a discussion. see. Ονεν Levensonetal. 2006)).," These more sensitive X-ray measurements would enable a test of the Compton-thick scenario, since better spectral measurements in the soft X-ray band could provide a possible measurement of the Fe K line, which is seen in heavily absorbed AGN (for a discussion, see, e.g., \citealt{LE06.1}) )."
 Seusitive hard N-rav measurements would provide iuch better constraiuts ou the absorbing column., Sensitive hard X-ray measurements would provide much better constraints on the absorbing column.
Telescope uarrow-band mesij of the eniss3on Could also show whethereit has thebiconical10LJ morphology expected for AGN , narrow-band imaging of the emission could also show whether it has the biconical morphology expected for AGN jets.
We report observations of a double X-ray. source with 1.9 De) kpe. or 0105. projected spatial separation in the :=01569 candidate dual AGN ealaxy600835.," We report observations of a double X-ray source with 1.9 $h^{-1}_{70}$ kpc, or $\farcs$ 68, projected spatial separation in the $z=0.1569$ candidate dual AGN galaxy."
"7.. This Sevtert 2 galaxy exhibits double-peaked emission lines with 350 lan s| line-of-sight1ΟLU) velocity separation in its SDSS spectruni and follow-up Lick/Ixast loneslit spectra show two 1.9 he,y onrhilekpe separation cussion colpoucuts."," This Seyfert 2 galaxy exhibits double-peaked emission lines with 350 km $^{-1}$ line-of-sight velocity separation in its SDSS spectrum, and our follow-up Lick/Kast longslit spectra show two 1.9 $h^{-1}_{70}$ kpc separation emission components."
" Ww the velocity aud spatial offsets provide circumstantial evidence for dual ACN. these features could also be produced by eas Lkineniatics from a sinele ACN,"," While the velocity and spatial offsets provide circumstantial evidence for dual AGN, these features could also be produced by gas kinematics from a single AGN."
 The observations bolster the evidence for dual AGN. by revealing two Nav components suggestive of Compton-thick ACN with the sanie spatial separation and oricutation as the two sources of optical eniission.," The observations bolster the evidence for dual AGN, by revealing two X-ray components suggestive of Compton-thick AGN with the same spatial separation and orientation as the two sources of optical emission."
 To date. dual AGN have typically been identified serendipitously because of the interesting characteristics of their host galaxies.," To date, dual AGN have typically been identified serendipitously because of the interesting characteristics of their host galaxies."
 These host salaxies include ultraluniunous mfrared galaxies2008).. a double radio source at the center of a galaxy cluster (Hudsouetal.2006).. aud a host ealaxy with double bright nuclei aud a tidal tail (Comerfordetal.2009b)..," These host galaxies include ultraluminous infrared galaxies, a double radio source at the center of a galaxy cluster \citep{HU06.1}, and a host galaxy with double bright nuclei and a tidal tail \citep{CO09.3}."
 is uulike these svsteuis because there is nothing particularly noteworthy about the galaxy. which is a secminely ordinary red sequence galaxy without tidal features visible in SDSS nuagiug or substructure reported in adaptive optics imaging.," is unlike these systems because there is nothing particularly noteworthy about the galaxy, which is a seemingly ordinary red sequence galaxy without tidal features visible in SDSS imaging or substructure reported in adaptive optics imaging."
 Our observations sugecst that dual AGN Ίαν be mere ubiquitous and not limited to ouly galaxies with extreme star formation or unusual morphologics., Our observations suggest that dual AGN may be more ubiquitous and not limited to only galaxies with extreme star formation or unusual morphologies.
 We have introduced a systematic. observational method for selecting promising dual ACN candidates. which have until now have ouly been ideutified through sereudipitous discoveries of individual svstenis.," We have introduced a systematic, observational method for selecting promising dual AGN candidates, which have until now have only been identified through serendipitous discoveries of individual systems."
 The method consists of three steps: 1) select dual AGN candidates as objects whose spectra exhibit double-peaked AGN ciuission lines in SDSS or other spectroscopic surveys of galaxies 2) conduct follow-up oueslit spectroscopy of the dual ACN candidates: 3) if the follow-up loneslit spectra reveal iu object has wo spatially-distinct ACN cunission coniponeuts. use ollow-up X-ray or radio observatious to identify whether he object is a dual AGN.," The method consists of three steps: 1) select dual AGN candidates as objects whose spectra exhibit double-peaked AGN emission lines in SDSS or other spectroscopic surveys of galaxies; 2) conduct follow-up longslit spectroscopy of the dual AGN candidates; 3) if the follow-up longslit spectra reveal an object has two spatially-distinct AGN emission components, use follow-up X-ray or radio observations to identify whether the object is a dual AGN."
 is the first object for which this techuique has been demonstrated. aud our observations show it most likely j)osts dual ACN: deeper XNorav observations would xovide the definitive evidence.," is the first object for which this technique has been demonstrated, and our observations show it most likely hosts dual AGN; deeper X-ray observations would provide the definitive evidence."
 Future observations will determine the eeneral applicability of this svstcmatic nethod for selecting dual ACN., Future observations will determine the general applicability of this systematic method for selecting dual AGN.
 JALC. acknowledecs insightfil discussions with Jeuux Greene. as well as support from a W.J. McDonald Postdoctoral Fellowship.," J.M.C. acknowledges insightful discussions with Jenny Greene, as well as support from a W.J. McDonald Postdoctoral Fellowship."
 The Texas Cosmology Center is supported by the College of Natural Sciences aud the Departiieunt. of Astronomy at the University of Texas at Austin and the MeDonald. Observatory., The Texas Cosmology Center is supported by the College of Natural Sciences and the Department of Astronomy at the University of Texas at Austin and the McDonald Observatory.
 B.F.C. aud C.NLM. were supported bv the U.S. ME of Euerev uuder coutract umber DE-ACO2-76SF0051, B.F.G. and G.M.M. were supported by the U.S. Department of Energy under contract number DE-AC02-76SF00515.
 , 
Models and simulations of jets produced by rotating magnetic fields generally assume the existence of an ordered. axially symmetric large-scale field of uniform polarity anchored in the central engine. starting with the original models of jets from accretion disks by ? and 2.. or the magnetic supernova model of ?..,"Models and simulations of jets produced by rotating magnetic fields generally assume the existence of an ordered, axially symmetric large-scale field of uniform polarity anchored in the central engine, starting with the original models of jets from accretion disks by \citet{1976Bisnovatyi} and \citet{1976Blandford}, or the magnetic supernova model of \citet{1970LeBlanc}."
 While such ordered fields are the most effective in producing Jets. the question whether they actually exist in accretion disks or the core of a star is still quite open.," While such ordered fields are the most effective in producing jets, the question whether they actually exist in accretion disks or the core of a star is still quite open."
 The largest scale at which magnetorotational (MRI) turbulence in accretion disks shapes the magnetic field structure is set by the disk thickness. which in. turn is much smaller than jets.," The largest scale at which magnetorotational (MRI) turbulence in accretion disks shapes the magnetic field structure is set by the disk thickness, which in turn is much smaller than jets."
 Collapsar cores are also small compared to the expected GRB jet. and it is not at all clear why the magnetic fields there should be ordered and axisymmetric.," Collapsar cores are also small compared to the expected GRB jet, and it is not at all clear why the magnetic fields there should be ordered and axisymmetric."
 Large-scale fields are not easily trapped by an accretion disk (?).., Large-scale fields are not easily trapped by an accretion disk \citep{1989Ballegooijen}.
 Without a large-scale field protruding from the disk. one may still hope to launch outflows by twisting fields inside the disk or by magnetic loops that extend into the disk corona as e.g. in 222.," Without a large-scale field protruding from the disk, one may still hope to launch outflows by twisting fields inside the disk or by magnetic loops that extend into the disk corona as e.g. in \citet{1979Galeev,1996Tout,2008Uzdensky}."
 The huge range of involved length scales is à major issue in Jet modeling. as numerical simulations that cover all scales are not feasible to date.," The huge range of involved length scales is a major issue in jet modeling, as numerical simulations that cover all scales are not feasible to date."
 Protostellar jets may be several parsecs long. launched by disks with sizes of «100 AU. re. about a factor 10 smaller (e.g.2)..," Protostellar jets may be several parsecs long, launched by disks with sizes of $\simm 100\,\mathrm{AU}$ , i.e. about a factor $10^4$ smaller \citep[e.g.][]{2001Shepherd}."
" Assuming that the ""launching scale"" is much smaller than the disk. perhaps on the order of -I1AU. one obtains an even bigger contrast in. length scales."," Assuming that the “launching scale” is much smaller than the disk, perhaps on the order of $\simm 1\,\mathrm{AU}$, one obtains an even bigger contrast in length scales."
 Supposing that jets in AGN are launched at a few Schwarzschild radit from the central black hole and taking Cygnus A as an example. one obtains a ratio of 10° between the jet length and the size of the engine (??)..," Supposing that jets in AGN are launched at a few Schwarzschild radii from the central black hole and taking Cygnus A as an example, one obtains a ratio of $10^6$ between the jet length and the size of the engine \citep{1998Krichbaum,2003Tadhunter}."
 The final jet properties are determined before it becomes ballistic. probably at scales which are somewhat smaller than that of the largest visible structures. but still considerably larger than the central engine.," The final jet properties are determined before it becomes ballistic, probably at scales which are somewhat smaller than that of the largest visible structures, but still considerably larger than the central engine."
 There is some numerical evidence that small-scale fields can also be used to generate jets., There is some numerical evidence that small-scale fields can also be used to generate jets.
 Most of these simulations are promising in terms of outflow production but limited to the immediate surroundings of the outflow-forming disk., Most of these simulations are promising in terms of outflow production but limited to the immediate surroundings of the outflow-forming disk.
 Axisymmetric simulations of outflows generated with small magnetic loops were done by ???..," Axisymmetric simulations of outflows generated with small magnetic loops were done by \citet{1998Romanova,1999Turner,2002Kudoh}."
" In their 3D simulations of accretion flows. ? demonstrated that an initially poloidal magnetic field confined within a rotating torus surrounding the acereting black hole can give rise to a transient outflow driven by accumulated toroidal fields in the form of a “magnetic tower"" (2)."," In their 3D simulations of accretion flows, \citet{2004Kato} demonstrated that an initially poloidal magnetic field confined within a rotating torus surrounding the accreting black hole can give rise to a transient outflow driven by accumulated toroidal fields in the form of a “magnetic tower” \citep{2003Lynden}."
 ? showed in simulations that loops of poloidal field in an accreting torus may give rise to a large-scale poloidal field as the field lines are stretched out in an axial outflow., \citet{2005DeVilliers} showed in simulations that loops of poloidal field in an accreting torus may give rise to a large-scale poloidal field as the field lines are stretched out in an axial outflow.
 The inflation and disruption of a magnetic loop outside a disk. caused by the generation of a toroidal field through differential rotation. was observed by ? in simulations of outflows from star-disk magnetospheres.," The inflation and disruption of a magnetic loop outside a disk, caused by the generation of a toroidal field through differential rotation, was observed by \citet{2009Fendt} in simulations of outflows from star-disk magnetospheres."
 The generation of magnetic flows within unmagnetized surroundings has also been studied in laboratory experiments (??):: such Jets are strongly affected by current-driven instabilities.," The generation of magnetic flows within unmagnetized surroundings has also been studied in laboratory experiments \citep{2005Hsu,2009Ciardi}; such jets are strongly affected by current-driven instabilities."
 The aim of the calculations presented here is to see whether small-scale fields in the form of loops anchored in a rotating disk can be used to produce jets of significant length (compared to the size of the source)., The aim of the calculations presented here is to see whether small-scale fields in the form of loops anchored in a rotating disk can be used to produce jets of significant length (compared to the size of the source).
 In the cases studied the flows propagate and are confined in an external unmagnetized atmosphere (as opposed to ? and ?.. hereafter Papers I&II. where we studied jets embedded in a large scale magnetically dominated environment).," In the cases studied the flows propagate and are confined in an external unmagnetized atmosphere (as opposed to \citealt{2008Moll} and \citealt{2009Moll}, hereafter Papers II, where we studied jets embedded in a large scale magnetically dominated environment)."
 While still idealized. this addresses environments like. protostellar jets launched into a dense cloud. or GRB jets launched by a collapsar core.," While still idealized, this addresses environments like protostellar jets launched into a dense cloud, or GRB jets launched by a collapsar core."
 To be investigated here are the circumstances under which jetlike flows are formed that penetrate through the atmosphere instead of dissipating in it., To be investigated here are the circumstances under which jetlike flows are formed that penetrate through the atmosphere instead of dissipating in it.
 As this is also a question of (non-axisymmetric) stability. three-dimensional simulations are necessary.," As this is also a question of (non-axisymmetric) stability, three-dimensional simulations are necessary."
 One of the questions to be answered is whether models of jets from small-scale magnetic fields are a viable alternative to those based on the twist of large-scale fields., One of the questions to be answered is whether models of jets from small-scale magnetic fields are a viable alternative to those based on the twist of large-scale fields.
 The primary ingredients in all our models are a rotating disk. implemented as a boundary condition. and magnetic field loops anchored in and sticking out of the disk.," The primary ingredients in all our models are a rotating disk, implemented as a boundary condition, and magnetic field loops anchored in and sticking out of the disk."
 The magnetic field geometries considered are sketched in Fig. 1.., The magnetic field geometries considered are sketched in Fig. \ref{fig:emergecases}.
 In case (a). some of the loops have one footpoint inside the disk while the other is anchored outside.," In case (a), some of the loops have one footpoint inside the disk while the other is anchored outside."
 The resulting shear motion of the footpoints generates a toroidal magnetic field., The resulting shear motion of the footpoints generates a toroidal magnetic field.
 In case (b). the loops are arranged such that they are not sheared.," In case (b), the loops are arranged such that they are not sheared."
 Here. a toroidal field can only beproduced by the inertia of the material above the," Here, a toroidal field can only beproduced by the inertia of the material above the"
Observations were made using 2 m telescope of IUCAA. Pune using the polarimeter (IMPOL) attached to it.,"Observations were made using 2 m telescope of IUCAA, Pune using the polarimeter (IMPOL) attached to it."
 The polarimeter available at the cassegrain focus of 2 m IUCAA telescope has a rotating hall-wave plate (IWDP) and Wollaston prism through which light passes belore forming a pair of images of an object on the CCD (Ramprakashοἱal.1993)., The polarimeter available at the cassegrain focus of 2 m IUCAA telescope has a rotating half-wave plate (HWP) and Wollaston prism through which light passes before forming a pair of images of an object on the CCD \citep{ramprakash98}.
. The ΠΗΝΙΟ can rotate in several discrete steps. such that its fast axis makes angles (a) with some relerence direction (generally celestial north-south).," The HWP can rotate in several discrete steps, such that its fast axis makes angles $\alpha$ ) with some reference direction (generally celestial north-south)."
 The light which is transmitted out of the Wollaston prism forms (wo images of anv celestial source on the CCD. with the ordinary. anc extraordinary set of ravs.," The light which is transmitted out of the Wollaston prism forms two images of any celestial source on the CCD, with the ordinary and extraordinary set of rays."
" For each object observations were taken at [our different positions of the IIWDP: QU. 229.5, 45"" and ο)..."," For each object observations were taken at four different positions of the HWP: $^{0}$ , $^{0}$ .5, $^{0}$ and $^{0}$ .5."
 Among BVRI polarimetry. the V. band is known to show the maximum polarization (Serkowskietal.1975). and is ideal for the detection of polarization caused due to scattering by dust.," Among BVRI polarimetry, the $V$ band is known to show the maximum polarization \citep{serkowski75} and is ideal for the detection of polarization caused due to scattering by dust."
 Moreover. because ofa low photon thax in the D band. the polarization measurements in (he B band often have large errors.," Moreover, because of a low photon flux in the $B$ band, the polarization measurements in the $B$ band often have large errors."
 For these reasons. we have conducted only. V-band polarimetry for the program stars.," For these reasons, we have conducted only $V$ -band polarimetry for the program stars."
 In addition to the program stars we have observed three polarization standard stars and one unpolarized standard star [ον polarimetric calibration., In addition to the program stars we have observed three polarization standard stars and one unpolarized standard star for polarimetric calibration.
 Data reduction was earried out using various tasks in ΗΑΕ., Data reduction was carried out using various tasks in IRAF.
" The task PILOT in APPILOT was used to measure the stellar Πας,", The task PHOT in APPHOT was used to measure the stellar flux.
 Zero polarization standard stars were observed to check for any possible instrumental error which proved to be smaller than 0.1%., Zero polarization standard stars were observed to check for any possible instrumental error which proved to be smaller than 0.1.
. For stus which show low deeree of polarization. measurement accuracy is an important issue.," For stars which show low degree of polarization, measurement accuracy is an important issue."
 The measured errors ave smaller than 0.1 for majority of the observed stars., The measured errors are smaller than 0.1 for majority of the observed stars.
 V-band polarimetric measurements for (he standard polarized ancl unpolarized stars are presented in Table 2., $V$ -band polarimetric measurements for the standard polarized and unpolarized stars are presented in Table 2.
 Estimates from literature are also listed for a comparison., Estimates from literature are also listed for a comparison.
 A close agreement between the estimated and literature values lends support to the reliability of our results., A close agreement between the estimated and literature values lends support to the reliability of our results.
 The estimated error in polarization is less than or near about for all stars (Table 3) except [ον the star WD 209621 for which the error is c 0.23., The estimated error in polarization is less than or near about for all stars (Table 3) except for the star HD 209621 for which the error is ${\pm}$ 0.23.
 Since polarimetric error in our case is photon noise dominated. in the case of very bright stars (his error becomes close {ο zero.," Since polarimetric error in our case is photon noise dominated, in the case of very bright stars this error becomes close to zero."
 Our program stus include objects brighter as well as fainter than WD 209621., Our program stars include objects brighter as well as fainter than HD 209621.
 Thus. the error in HD 209621 can be considered (o be a tvpical error value.," Thus, the error in HD 209621 can be considered to be a typical error value."
" The polarization standard stars LID 147084 and ID 160529. both observed on the same night. show an offset of about &—10"": the near equality of the olfsets is as expected Lor consistency. and a similaroffset was also reported earlier bv Senetal.(2000).."," The polarization standard stars HD 147084 and HD 160529, both observed on the same night, show an offset of about ${\sim}-$ $^{0}$; the near equality of the offsets is as expected for consistency, and a similaroffset was also reported earlier by \citet{sen00}."
 The offset is, The offset is
"L6pt PACS muubers: i. j. 95.55.Ry. 3.85ΕΤ). 98.70.Sa Iu convention. it is convinced that the dominant energy loss is PP instead of ES in the Standard Model (SM) for the cosimic ray photon with cucrey above the PP threshold Ly, [1] where e is the energy of à background photon.","16pt PACS numbers: $-$ i, $-$ j, 95.85.Ry, 13.85.Tp, 98.70.Sa In convention, it is convinced that the dominant energy loss is PP instead of ES in the Standard Model (SM) for the cosmic ray photon with energy above the PP threshold $E_{\rm th}$ \cite{review-ray}
 where $\epsilon$ is the energy of a background photon."
 However recent research [2.3.1] ou diphoton interaction reveals that the cross section of uuparticle exchange can easily surpass the SM one at high enough energy because unparticle exchanges are also at the tree-level through all s-. t+. aud e-chauncls.," However recent research \cite{EC1,EC2,EC3} on diphoton interaction reveals that the cross section of unparticle exchange can easily surpass the SM one at high enough energy because unparticle exchanges are also at the tree-level through all $s$ -, $t$ -, and $u$ -channels."
 It is natural to the of plavsics ou thle cosmic ray. especially ou whether the appearanceexplore of consequencewill lead to its uuparticledominant energy loss process to photon.change from PP to ES. which will cause nontrivial uuparticleobservational in the spectrum of cosiic rav photon.," It is natural to explore the consequence of unparticle physics on the cosmic ray photon, especially on whether the appearance of unparticle will lead to its dominant energy loss process to change from PP to ES, which will cause nontrivial observational signals in the spectrum of cosmic ray photon."
" Iu the meamvhile. very receutlv the Pamela collaborationsignals announced their first measurements on the cosmic ray (CR) positron fraction [5] in the energy range 1.5. 100€6V. The positron fraction of Pamela data shows a promincut excess to the background estimation |6.7| of the conventional CR propagation model iu the region 10 00ο, This result is cousisteut with previous measuremoeuts by. eg. HEAT [8] aud AMS [9].."," In the meanwhile, very recently the Pamela collaboration announced their first measurements on the cosmic ray (CR) positron fraction \cite{p1} in the energy range $1.5-100$ GeV. The positron fraction of Pamela data shows a prominent excess to the background estimation \cite{t1,t2} of the conventional CR propagation model in the region $\sim10-100$ GeV. This result is consistent with previous measurements by, e.g., HEAT \cite{p3} and AMS \cite{p4}."
 On the other haud. the electron spectrum up to several TeV imieasured by ATIC collaboration also displays an obvious excess in the region around 300~ sO0CeV. [LO]. which confirms the measurements of the electron spectitun by PPB-BETS ΠΠ. I.E.S.5. [12.13].. aud most recently by Fermi |11]..," On the other hand, the electron spectrum up to several TeV measured by ATIC collaboration also displays an obvious excess in the region around $300 \sim 800 $ GeV \cite{e5}, which confirms the measurements of the electron spectrum by PPB-BETS \cite{e6}, H.E.S.S. \cite{e7,e8}, and most recently by Fermi \cite{e9}."
" The mismatch between theory and observatious stimulates a lot of interest ou the cosmic rav ¢+) and we will reexamine the propagation of cosmic rav ο in the framework of unparticle physics,"," The mismatch between theory and observations stimulates a lot of interest on the cosmic ray $e^\pm$, and we will reexamine the propagation of cosmic ray $e^\pm$ in the framework of unparticle physics."
 Ou particular. we address the dominant loss process for cose rav 67. inverse Compton scattering. to study the impact of wuparticle stuff on c and further ou the observational CXCCSS.," On particular, we address the dominant loss process for cosmic ray $e^\pm$, inverse Compton scattering, to study the impact of unparticle stuff on $e^\pm$ and further on the observational excess."
The is as follows., The paper is organized as follows.
 Iu the next section. we overview the basic of uuparticle paper the organizedodd and space of ποιο stuff with propertydifferent Lorentz physics.structures.," In the next section, we overview the basic property of unparticle physics, including the odd propagator and phase space of unparticle stuff with different Lorentz structures."
 In inclidingSec.3. we derive the propagatorscattering phaseamplitudes for the πηρανinvolved processes. that is. the PP," In Sec.3, we derive the scattering amplitudes for the involved processes, that is, the PP"
 inthe. 9 ΙΟHard Xrayobsercationswillth usheiuiportait., in the 0.1-2.4 keV. Hard X-ray observations will thus be important.
 As discussed in SectionLL. the compact hot dust componcut that we found in LEDÀ 81271 and IRAS 01250|2832 is likely to surround an ACN that is heating it. even thoueh other possibilities cannot be completely ruled out.," As discussed in Section\ref{sec:hot}, the compact hot dust component that we found in LEDA 84274 and IRAS 01250+2832 is likely to surround an AGN that is heating it, even though other possibilities cannot be completely ruled out."
 Iu this section. we discuss these dusty ACNs.," In this section, we discuss these dusty AGNs."
 The optical spectra display no strongC» emission lines (especially hieh-excitation[m] ones) from nurow-line regions., The optical spectra display no strong emission lines (especially high-excitation ones) from narrow-line regions.
 Thus. no nirrow-liue region ou a 10-1000pc scale is evideut in these galaxies.," Thus, no narrow-line region on a 10-1000pc scale is evident in these galaxies."
 IRAS 01250|2832 has a low-ionization spectrun that is tvpical of LLAGNs. as discussed in Section 3.2..," IRAS 01250+2832 has a low-ionization spectrum that is typical of LLAGNs, as discussed in Section \ref{sec:iras}."
 The low-ionization state spectrum in LLACGNs may be produced by black-hole accretion at a very low rate (with a bolometric Iuninositv of Lago<Ls1019Le 2008)))., The low-ionization state spectrum in LLAGNs may be produced by black-hole accretion at a very low rate (with a bolometric luminosity of $L_{bolo}<4\times 10^{10}\ \mathrm{L}_{\sun}$ ).
 Towever. LEDA 81271 and IRAS 01250|2832 have much higher luminosities of Lig21«1014Ls at infrared waveleneths than LLAGNs.," However, LEDA 84274 and IRAS 01250+2832 have much higher luminosities of $L_{\mathrm{IR}}
\gtrsim 1 \times 10^{11}\ \mathrm{L}_{\sun}$ at infrared wavelengths than LLAGNs."
 Large optical extinction of the central eugiue is a plausible explanation of their optical spectra., Large optical extinction of the central engine is a plausible explanation of their optical spectra.
 These findings indicate that the AGNs are surrounded by a larec amount of dust. so that ionizing UV radiation from the AGN is blocked along virtually all lines of sight.," These findings indicate that the AGNs are surrounded by a large amount of dust, so that ionizing UV radiation from the AGN is blocked along virtually all lines of sight."
 Iu ultra-huninous aud bhbuumous infrared ealaxy cores. similar situations are ποιολος seen2010).," In ultra-luminous and luminous infrared galaxy cores, similar situations are sometimes seen."
". Since the observed radiation from ACUNS. is often reprocessed and re-enmütted at infrared waveleusths. the iutriusie SED is altered subsautially,"," Since the observed radiation from AGNs is often reprocessed and re-emitted at infrared wavelengths, the intrinsic SED is altered substantially."
" Wide spectral coverage Is necessary to determine the Eddington ratio (Li4,/LEdd» Or equivalently Lyon/ Abu)1997)."," Wide spectral coverage is necessary to determine the Eddington ratio $L_{\mathrm{bolo}}/L_{\mathrm{Edd}}$, or equivalently $L_{\mathrm{bolo}}/M_{\mathrm{BH}}$ )."
. Although it is difficult to differeifiate t16 ACN component fou the overall SED. we estimate the black-hole mass by naking the a simple assuniption that the ACN luminosity is nearly equal to the luninosity roni the 500-GOO0IN. dust.," Although it is difficult to differentiate the AGN component from the overall SED, we estimate the black-hole mass by making the a simple assumption that the AGN luminosity is nearly equal to the luminosity from the 500-600K dust."
 Asstuuing spherical accretion toward the black hole wit ja radiative efficiency of 01. the supermassive black hole mass from the Iuniuosities oft1ο hot cust is at least 8«1094 and 2«107M. for LEDA 8127 Laud IRAS 01250|2832. respectively.," Assuming spherical accretion toward the black hole with a radiative efficiency of 0.1, the supermassive black hole mass from the luminosities of the hot dust is at least $8\times 10^6\ 
\mathrm{M}_{\odot}$ and $2\times 10^7\ \mathrm{M}_{\odot}$ for LEDA 84274 and IRAS 01250+2832, respectively."
 These do not differ siguificautly from the black hole masses estimated from the luminosities using the volometric correction evaluated using a Type2 QSO teiiplate in«, These do not differ significantly from the black hole masses estimated from the luminosities using the bolometric correction evaluated using a Type2 QSO template in.
 Plus. we adopt the black hole mass estimates from the luninositics of the hot αιραε in the following discussion.," Thus, we adopt the black hole mass estimates from the luminosities of the hot dust in the following discussion."
 Ifthey have a higher radiative efficiency. the masses will be OWOL.," If they have a higher radiative efficiency, the masses will be lower."
 The masses of these host galaxies can be calculated to be ος107AL. aud £s10?N or LEDA 81271 «ud IRAS 01250|2832. respectively. using the 2MASS photometry," The masses of these host galaxies can be calculated to be $6\times 10^9\ \mathrm{M}_{\odot}$ and $4
\times 10^{9}\ \mathrm{M}_{\odot}$ for LEDA 84274 and IRAS 01250+2832, respectively, using the 2MASS photometry"
The specificity of the Sun and of our solar system have been the subject of active investigation over the last 5 decades.,The specificity of the Sun and of our solar system have been the subject of active investigation over the last 5 decades.
 How typical is the Sun for à star of its age. mass. and chemical composition?," How typical is the Sun for a star of its age, mass, and chemical composition?"
 How typical is that solar-type stars host planetary systems?, How typical is that solar-type stars host planetary systems?
 Are they similar at all to ours?, Are they similar at all to ours?
 The quest to find stellar analogues to the Sun has been going on for a long time (for an extensive review see. e.g.. CayreldeStrobel 1996)). and it stems from the poor knowledge we have of the Sun when seen “as a star’ and from how typical the Sun is for a G2 type star. for its age. chemical composition. population.," The quest to find stellar analogues to the Sun has been going on for a long time (for an extensive review see, e.g., \citealt{Cayrel1996}) ), and it stems from the poor knowledge we have of the Sun when seen `as a star' and from how typical the Sun is for a G2 type star, for its age, chemical composition, population."
 It is. however. after the discovery of the first exo-planets (Mayor&Queloz 1995)) that this quest became even more compelling. because to find stars similar to our own would allow us to answer to fundamental questions related to the origin of the solar system. the frequency of planetary systems similar to ours. and eventually the formation of life in other exo-planetary systems (CayreldeStrobel 1996)).," It is, however, after the discovery of the first exo-planets \citealt{MayQue1995}) ) that this quest became even more compelling, because to find stars similar to our own would allow us to answer to fundamental questions related to the origin of the solar system, the frequency of planetary systems similar to ours, and eventually the formation of life in other exo-planetary systems \citealt{Cayrel1996}) )."
 The need to identify in the night sky solar proxies to be used for spectroscopic comparison is also diffuse. in particular for the analysis of small solar system bodies (Bóhhnhardt. private communication).," The need to identify in the night sky solar proxies to be used for spectroscopic comparison is also diffuse, in particular for the analysis of small solar system bodies (Böhhnhardt, private communication)."
 Among the most recent results in this research. Meléndezetal.(2006) used high resolution. high signal-to-noise ratio Keck spectra to show that HD 98618 is a very close solar twin. and Kingetal.(2005) proposed HD 143436 after analyzing 4 stars pre-selected from literature.," Among the most recent results in this research, \cite{Melendez2006} used high resolution, high signal-to-noise ratio Keck spectra to show that HD 98618 is a very close solar twin, and \cite{King2005} proposed HD 143436 after analyzing 4 stars pre-selected from literature."
 These stars seem to compare well with the best known solar twin. HR 6060. first analyzed by PortodeMello&daSilva(1997).. and subsequently confirmed by Soubiran&Triaud(2004).. who made a comparative study of several hundreds of ELODIE spectra.," These stars seem to compare well with the best known solar twin, HR 6060, first analyzed by \cite{PorSil1997}, and subsequently confirmed by \cite{SoubTri2004}, who made a comparative study of several hundreds of ELODIE spectra."
 Finally. Meléndez&Ramírez(2007) have shown HIP 56948 to be the best solar twin known to date both in stellar parameters and in chemical composition. including a low lithium abundance.," Finally, \cite{Melendez2007} have shown HIP 56948 to be the best solar twin known to date both in stellar parameters and in chemical composition, including a low lithium abundance."
 The open cluster M67 is a perfect target to search for solar analogues., The open cluster M67 is a perfect target to search for solar analogues.
 Recent chemical analyses (TautvaiSieneetal.2000;Randichetal.2006:Pace 2008)). show that this cluster has a chemical composition (not only Fe. but also all the other elements) extremely similar to the solar one. as close as allowed by the high precision of the measurements.," Recent chemical analyses \citealt{Tau2000, Randich2006, Pace2008}) ), show that this cluster has a chemical composition (not only Fe, but also all the other elements) extremely similar to the solar one, as close as allowed by the high precision of the measurements."
 The analysis resulted in. [Fe/H]=—-0.03+40.03 for Tautvaisieneetal.(2000)... [Fe/H]=0.03+40.01 for Randichetal.(2006).. and Fe/H]=0.03+0.03 for Paceetal.(2008).," The analysis resulted in $-$ $\pm$ 0.03 for \cite{Tau2000}, $\pm$ 0.01 for \cite{Randich2006}, and $\pm$ 0.03 for \cite{Pace2008}."
. There are other two additional characteristics Which make M67 strategical., There are other two additional characteristics which make M67 strategical.
 The first one is that all the determinations of age give for this cluster an age encompassing that of the Sun (3.5-4.8 Gyr; Yadavetal. 2008)). while the age determination for field stars is always uncertain.," The first one is that all the determinations of age give for this cluster an age encompassing that of the Sun (3.5-4.8 Gyr; \citealt{Yadav2008}) ), while the age determination for field stars is always uncertain."
 The second characteristic 15 that M67 is among the very few clusters showing Li depleted G stars (Pasquinietal. 1997))., The second characteristic is that M67 is among the very few clusters showing Li depleted G stars \citealt{Pasquini1997}) ).
 This is an important point because. as pointed out by CayreldeStrobel(1996).. even if many stars appear to have most characteristies similar to the Sun. their Li abundance is usually 10 times higher than in ot star.," This is an important point because, as pointed out by \cite{Cayrel1996}, even if many stars appear to have most characteristics similar to the Sun, their Li abundance is usually 10 times higher than in our star."
 Since Li is likely a indicator of the complex interaction taking place in the past between the stellar external layers and the hotter interior. the choice of stars which also share the same Li abundance with the Sun is an additional property to pinpoint the true analogues.," Since Li is likely an indicator of the complex interaction taking place in the past between the stellar external layers and the hotter interior, the choice of stars which also share the same Li abundance with the Sun is an additional property to pinpoint the true analogues."
 In our opinion. the search of analogues to the Sun and to the solar system can be well performed in open clusters (OCs). which show a homogeneous age and chemical composition. common birth and early dynamical environment.," In our opinion, the search of analogues to the Sun and to the solar system can be well performed in open clusters (OCs), which show a homogeneous age and chemical composition, common birth and early dynamical environment."
 As à consequence. they provide an excellent laboratory for investigating the physics of solar stars and of planetary system evolution. besides being excellent probes of the structure and evolution of the Galactic disk.," As a consequence, they provide an excellent laboratory for investigating the physics of solar stars and of planetary system evolution, besides being excellent probes of the structure and evolution of the Galactic disk."
 M67 is a rich. cluster. therefore it provides us with the opportunity to find many stars candidates sharing. similar characteristics. and not only one.," M67 is a rich cluster, therefore it provides us with the opportunity to find many stars candidates sharing similar characteristics, and not only one."
 This ts fundamental to obtain some meaningful statistics. and the cluster hosts many main sequence (MS) stars of mass around the solar mass. which form a continuous distribution (Fig. 1)).," This is fundamental to obtain some meaningful statistics, and the cluster hosts many main sequence (MS) stars of mass around the solar mass, which form a continuous distribution (Fig. \ref{fig:cmd_M67}) )."
"The multiphase nature of the intracluster medium (1031) in cooling Hows was demonstrated. a decade ago when deprojections of X-ray surface-brightness profiles showed that mass cools ancl is. deposited. from. the flow in a distributed manner. AXr ""Thomas. Fabian Nulsen LOST).","The multiphase nature of the intracluster medium ) in cooling flows was demonstrated a decade ago when deprojections of X-ray surface-brightness profiles showed that mass cools and is deposited from the flow in a distributed manner, $\dot{M}\approxpropto r$ Thomas, Fabian Nulsen 1987)."
 llowever. the complexity of the theory and lack of data with high spatia resolution means that the single-phase approximation is still widely adopted.," However, the complexity of the theory and lack of data with high spatial resolution means that the single-phase approximation is still widely adopted."
 In this paper E show how the use of self-similar cdensitv clistributions can [ead to a relatively simple form for the multiphase cooling How equations whils spanning the whole range of expected behaviours in the more general case., In this paper I show how the use of self-similar density distributions can lead to a relatively simple form for the multiphase cooling flow equations whilst spanning the whole range of expected behaviours in the more general case.
 In combination with the emissivity profile (for a spherically-svnimetric flow). the equations can be solvec o vield the gas temperature and mass profiles within the cooling Low.," In combination with the emissivity profile (for a spherically-symmetric flow), the equations can be solved to yield the gas temperature and mass profiles within the cooling flow."
 In Section 2.. L introduce the concept. of the volume raction to describe the distribution of density phases in heLOCAL.," In Section \ref{sec:dd}, I introduce the concept of the volume fraction to describe the distribution of density phases in the."
 The self-similar forms of the volume fraction are determined and shown to span all reasonable behaviours for he more general situation., The self-similar forms of the volume fraction are determined and shown to span all reasonable behaviours for the more general situation.
 The steady-state. self-similar orm of the cooling [low equations are derived in 3 and it is shown how these can be solved in the spherically-svnimetric case if the emissivity profile is known.," The steady-state, self-similar form of the cooling flow equations are derived in \ref{sec:theory} and it is shown how these can be solved in the spherically-symmetric case if the emissivity profile is known."
 The behaviour of the solutions is examined in detail: not every emissivity. profile is consistent withthese mocdels., The behaviour of the solutions is examined in detail: not every emissivity profile is consistent withthese models.
 In ει. the theory is applied to ROSANT. HIR data for three Abell clusters with some success.," In \ref{sec:abell}, the theory is applied to ROSAT HRI data for three Abell clusters with some success."
 Finally. the results are summarisecd ane further cliscussecl in 5..," Finally, the results are summarised and further discussed in \ref{sec:conc}."
 The theory of multiphase cooling Lows was set out by Nulsen (1986)., The theory of multiphase cooling flows was set out by Nulsen (1986).
 Hle introduced the concept of the volume fraction. v.[Gp0. such that fdp is the fractional volume occupied bv phases with densities in the range p to p|dp.," He introduced the concept of the volume fraction, $f(\rho,\rvec,t)$, such that $f\,\dd\rho$ is the fractional volume occupied by phases with densities in the range $\rho$ to $\rho+\dd\rho$."
" Dt is assumed that the phases comove but are thermally isolated from one anotherthese can be regarded as empirical facts as otherwise the How would rapidly evolve to a single-phase state,", It is assumed that the phases comove but are thermally isolated from one another—these can be regarded as empirical facts as otherwise the flow would rapidly evolve to a single-phase state.
 Given jese assumptions then it is relatively straightforward to derive an equation for the evolution of the volume fraction (sce Appendix)., Given these assumptions then it is relatively straightforward to derive an equation for the evolution of the volume fraction (see Appendix).
 Writing where and p=pir. 0. then," Writing ,t), where and $\rho=\rho(\rvec,t)$ , then"
density. gas-rich planets increases with planet size. a feature essentially intrinsic (o our analvsis because of our choice of a=2.,"density, gas-rich planets increases with planet size, a feature essentially intrinsic to our analysis because of our choice of $\alpha = 2$."
 Theoretical models predict the formation of inner. rocky planets (Ravimondοἱal.2004)... and the stellar UV-driven escape of any. primordial hydrogen atinosphleres (Piervehumbert&Gaidos2011)..," Theoretical models predict the formation of inner, rocky planets \citep{Raymond2004}, and the stellar UV-driven escape of any primordial hydrogen atmospheres \citep{Pierrehumbert2011a}."
 The low densitv of the short-period super-Earth GJ 1214b can be explained by a substantial hydrogen envelope Crolletal.201 1).. but also bv a thick IO shell 2011)..," The low density of the short-period super-Earth GJ 1214b can be explained by a substantial hydrogen envelope \citep{Charbonneau2009a,Croll2011}, but also by a thick $_2$ O shell \citep{Bean2010a,Desert2011}."
 Its host is a cooler (3000 IN). muehli less luminous nid-M star and (hiis may permil retention of hydrogen (Pierrehumbert&Gaidos2011)..," Its host is a cooler (3000 K), much less luminous mid-M star and this may permit retention of hydrogen \citep{Pierrehumbert2011a}."
 Gas- and ice-rich planets resembling GJ 1214b may be the exception rather than (he rule around (the coolest.Aepler Larget stars., Gas- and ice-rich planets resembling GJ 1214b may be the exception rather than the rule around the coolest target stars.
 Refinement of Doppler svstematic errors and the properties ofAepler target stars. specifically the radii of Ix. and M. dwarls aud the fraction of interloping giants. will perimit more robust constraints.," Refinement of Doppler systematic errors and the properties of target stars, specifically the radii of K and M dwarfs and the fraction of interloping giants, will permit more robust constraints."
 This research was supported by NSF erants AST-09-08406 CEG). AST-10-36283 (DAF). AST-09-08419 (SL): NASA grants NNXIOAI90CT and NNNILAC33G (EG). and the NASA KPDA program (DAF).," This research was supported by NSF grants AST-09-08406 (EG), AST-10-36283 (DAF), AST-09-08419 (SL); NASA grants NNX10AI90G and NNX11AC33G (EG), and the NASA KPDA program (DAF)."
 TheAepler mission is funded by the NASA Science Mission Directorate. and data were obtained from the Multimission Archive at the Space Telescope Science Institute. funded by NASA grant NNXNOSAFOSG. We thank Andrew Lowarel for his help with screening Doppler template spectra.," The mission is funded by the NASA Science Mission Directorate, and data were obtained from the Multimission Archive at the Space Telescope Science Institute, funded by NASA grant NNX09AF08G. We thank Andrew Howard for his help with screening Doppler template spectra."
iue of sight.,line of sight.
" It is. therefore. reasonable to suppose that what we lea‘tn about the physical state of his system. in particular its cisς, by exploiting the fortuitous geometry will ‘eliably inform studies oL protoplanetary disks iu geneal."," It is, therefore, reasonable to suppose that what we learn about the physical state of this system, in particular its disk, by exploiting the fortuitous geometry will reliably inform studies of protoplanetary disks in general."
 With a total mass near 1 M. ald an age of ~8 My. INH 15D uay be particularly relevant to he stage of planet formation in our own Sola“System curing which he most primitive meteorites aud their embedded choudrules formed.," With a total mass near 1 $_\odot$ and an age of $\sim$ 3 My, KH 15D may be particularly relevant to the stage of planet formation in our own Solar System during which the most primitive meteorites and their embedded chondrules formed."
 Plavelian have recently identified two other systems with NH 15D-like variations aud have proposed tliat hey form a new class of variable stars with warped iuner disks caused by the existence of tertiary components.," \citet{p08,p10} have recently identified two other systems with KH 15D-like variations and have proposed that they form a new class of variable stars with warped inner disks caused by the existence of tertiary components."
 There is possible evideuce lor third light in the photometry of WH 15D presented iu his paper. iu support of that idea.," There is possible evidence for third light in the photometry of KH 15D presented in this paper, in support of that idea."
 This paper presents photometric data in V and I over the period 2005-10 aucl au empirical description and basic analysis of those data., This paper presents photometric data in V and I over the period 2005-10 and an empirical description and basic analysis of those data.
 Π will hopefully be useful to a next. generation of dynamical aud. phenomenological moclels along the lines of those proposed by Chiang (2001).. Winnetal.(2001.2006) and Silvia&Agol(2008).," It will hopefully be useful to a next generation of dynamical and phenomenological models along the lines of those proposed by \citet{CM04}, \citet{w04,w06} and \citet{sa08}."
. It is clear that none of the current models accurately predicted the observed behavior reported here. nor do any of them incorporate all of the relevant physical components. πιο as a scattered light source with the observed properties.," It is clear that none of the current models accurately predicted the observed behavior reported here, nor do any of them incorporate all of the relevant physical components, such as a scattered light source with the observed properties."
 [t is likely that the next. generation of models will make it possible to exploit more fully. aud coulicently the information in the cata presented here than we can do empirically., It is likely that the next generation of models will make it possible to exploit more fully and confidently the information in the data presented here than we can do empirically.
 The current paper will thus be restricted to preseutiug the data auc providing some qualitative guidauce to its interpretation based ou tlie perspective of observers., The current paper will thus be restricted to presenting the data and providing some qualitative guidance to its interpretation based on the perspective of observers.
 Iu this paper we report V and I fou the Cousins system) photometric data on WH 15D obtaiued at the Vau Vleck Observatory (VVO ou the campus of Weslevan University in Middletown. C'T. at Cerro Tololo Inter-American Observatory (CTIO) in Chile with a SMABTS consortium telescope and from Mount Maidanalk Observaory (MIMO) in Uzbekistan. operated by the Uzbek Academy ol Scieuces.," In this paper we report V and I (on the Cousins system) photometric data on KH 15D obtained at the Van Vleck Observatory (VVO) on the campus of Wesleyan University in Middletown, CT, at Cerro Tololo Inter-American Observatory (CTIO) in Chile with a SMARTS consortium telescope and from Mount Maidanak Observatory (MMO) in Uzbekistan, operated by the Uzbek Academy of Sciences."
 Five vears of observations from September. 2005. through Mare. 2010. are preseutecl.," Five years of observations from September, 2005, through March, 2010, are presented."
 At VVO. data were obtained with the 0.6 im Perkin telescope.," At VVO, data were obtained with the 0.6 m Perkin telescope."
 In 2005-6 acL 2006-7 we employed a 1021 x 1021 [ront-illuminuated Photometrics CCD camera., In 2005-6 and 2006-7 we employed a 1024 x 1024 front-illuminated Photometrics CCD camera.
 Five successive one-minute exposures in I were taken ouce each night and suimuned., Five successive one-minute exposures in I were taken once each night and summed.
" Since 2007-8 the CCD camera as been a 2018 x 2018 back-illuuninated Apogee camera with 0.3"" pixels.", Since 2007-8 the CCD camera has been a 2048 x 2048 back-illuminated Apogee camera with $\arcsec$ pixels.
 Between five ancl filtee: consecutive l-o1ninute curation Lor V images were obtaiued ou clear nights with this system aud suuned., Between five and fifteen consecutive 1-minute duration I or V images were obtained on clear nights with this system and summed.
" Typical seeing at VVO is about 2.5"" and the sky is rather bright. so the daa are more precise by a [actor of 2-3 during the brighter portious of the system's rauge than during the fainter."," Typical seeing at VVO is about $\arcsec$ and the sky is rather bright, so the data are more precise by a factor of 2-3 during the brighter portions of the system's range than during the fainter."
 At CTIO. images were obtained each clear nieht in all five observing seasous from the begiuuiug oL October until the eud of March with a 2018 x 2018 Fairchik CCD attached to the 1.3 1n telescope of the SMARTS consortium.," At CTIO, images were obtained each clear night in all five observing seasons from the beginning of October until the end of March with a 2048 x 2048 Fairchild CCD attached to the 1.3 m telescope of the SMARTS consortium."
" All images employed 2 x 2 binning and had a 6.3 x 6.3"" field of view.", All images employed 2 x 2 binning and had a $\arcmin$ x $\arcmin$ field of view.
 Four images of 5 s each were taken while WH 15D was it the bright phase aud normally eight, Four images of 5 s each were taken while KH 15D was in the bright phase and normally eight
 (ultra-compactdwarfgalaxiesorUCDs:A\Gieske (Barnes&Ieruquist.1992).. (Cappellari, \citep[ultra-compact dwarf galaxies or UCDs;][]{mieske} \citep{barnes}. \citep{cappellari}.
"ctal.2007).. (Zaritsky Y, witlin the halt-light radius à is one gauge of the relative concentration of barvous to dark matter."," \citep{zgz, zgzb, zzg} $\Upsilon_e$ within the half-light radius $r_{e}$ is one gauge of the relative concentration of baryons to dark matter."
" The larec Ἐν Ἡ of dwarf spheroidals (dSplis) reflect. the importance of dark imatter within their lunünous region. while the ower Y,"" of ellipticals arise from the dominance of the xwvons there."," The large $\Upsilon_e$ 's of dwarf spheroidals (dSphs) reflect the importance of dark matter within their luminous region, while the lower $\Upsilon_e$ 's of ellipticals arise from the dominance of the baryons there."
 One could hwpothesize a more extreme case than that for cllipticals. where the barvous are yacked so tightly that the dark matter is dvuamically icelieible within +.," One could hypothesize a more extreme case than that for ellipticals, where the baryons are packed so tightly that the dark matter is dynamically negligible within $r_{e}$."
" It would be dithcult to differcutiate such a low Y, galaxy. if it were low mass. from a star cluster. which is devoid of dark imatter."," It would be difficult to differentiate such a low $\Upsilon_e$ galaxy, if it were low mass, from a star cluster, which is devoid of dark matter."
 If they inve anv dark matter. UCDs (Mieskeetal.2008) may (0 exanrples of galaxies whose spatial aud kincmatic structure thus resemble those of star clusters.," If they have any dark matter, UCDs \citep{mieske} may be examples of galaxies whose spatial and kinematic structure thus resemble those of star clusters."
 Are rere other observables that reveal their presuimablv distinct formation aud evolutionary paths?, Are there other observables that reveal their presumably distinct formation and evolutionary paths?
" Remarkably. colmparably low mass svstenis can have low Y, (stellar clusters. UCDs) or high Y. (dSplis)."," Remarkably, comparably low mass systems can have low $\Upsilon_e$ (stellar clusters, UCDs) or high $\Upsilon_e$ (dSphs)."
 What variations in dark halo properties of barvonic processes explain this wide range of barvon to dark matter concentrations?, What variations in dark halo properties of baryonic processes explain this wide range of baryon to dark matter concentrations?
 Such questions are fundamental to understanding the nature of low ass stellar svstenis. the basic characteristics of ealaxies. aud the wavs in which baryous can settle into a dark matter halo.," Such questions are fundamental to understanding the nature of low mass stellar systems, the basic characteristics of galaxies, and the ways in which baryons can settle into a dark matter halo."
 Here we use a new scaling relation (the Fundamental \lanitold of Zavitskvetal.(FAL:2006a. 2008))) to explore the first question above. whether star clusters can be distinguished. frou low mass galaxies with similar Y audtheimplications for their evolution.," Here we use a new scaling relation (the Fundamental Manifold of \citet[FM;][]{zgz,zzg}) ) to explore the first question above, whether star clusters can be distinguished from low mass galaxies with similar $\Upsilon_e$ 's andtheimplications for their evolution."
 Tn certain projectious ofparameter space. elobular clusters are distinct from galaxies," In certain projections ofparameter space, globular clusters are distinct from galaxies"
DAnegle 2 -,DAngle = .
 For a CTIL.ση)=1/1. and the inside lens equation is given bv ," For a CTIL,$\tilde\sigma(t) = 1/t$, and the inside lens equation is given by = z -."
For galaxy lenses (hat have elliptic shapes and light distributions. one can consider a class of mass distributions whose projected densities depend on only the elliptical radius.," For galaxy lenses that have elliptic shapes and light distributions, one can consider a class of mass distributions whose projected densities depend on only the elliptical radius."
 Forexample. M(/)x1// where the position variable £=/(esin+ibeos0).," Forexample, $\Sigma(t) \propto 1/t$ where the position variable $\xi = t(a\sin\theta + i b\cos\theta)$."
 It is currently unknown ifa singular elliptic mass distribution with infinite or finite extension is exactly or Closely related. {ο a sell-gravitating gas mass in thermal equilibrium., It is currently unknown if a singular elliptic mass distribution with infinite or finite extension is exactly or closely related to a self-gravitating gas mass in thermal equilibrium.
 Ilowever one can suspect that it might be the case at least for small elliplicilies because of the existence of the circularly svmmetric svstem SIS and the fact that an angular momentum tends to add oblateness to the svstem., However one can suspect that it might be the case at least for small ellipticities because of the existence of the circularly symmetric system SIS and the fact that an angular momentum tends to add oblateness to the system.
 Cousicer the family of ellipses parameterized by /(>0) that fills the (wo-climensional space., Consider the family of ellipses parameterized by $t ( \geq 0)$ that fills the two-dimensional space.
 — l(acos + ibsin lhela)).. a bI) The curve /=constant is an ellipse that is converted to the faaniliar equation Lorthe ellipse in real coordinates £4 and £o where £=€;+ /£.," = t (a + i b ), a b The curve $t = constant$ is an ellipse that is converted to the familiar equation forthe ellipse in real coordinates $\xi_1$ and $\xi_2$ where $\xi = \xi_1 + i\xi_2$ ."
," + = t^2 , t 0"
solid circle) [or dust reddening ancl the ratios of the adjacent spatial regions (evan symbols) are successfully explained by AGN photoionization along both PAs.,solid circle) for dust reddening and the ratios of the adjacent spatial regions (cyan symbols) are successfully explained by AGN photoionization along both PAs.
 On the contrary. the spectrum of the companion (nic2) shows very low line ratios. noticeably lower than those measured in typical vpe .2 AGNs.," On the contrary, the spectrum of the companion $nuc2$ ) shows very low line ratios, noticeably lower than those measured in typical type 2 AGNs."
wes Similar(quu values are often⋅ measured in LIL galaxies (e.g. Lamarcille et al. 2004))., Similar values are often measured in HII galaxies (e.g. Lamareille et al. \citeyear{lam04}) ).
 Dust reddening could shift an AGN to the LIL galaxy. region in the diagram. since the ΟΠΗ) would. be observed. to be lower. while /LLA7 would be similar.," Dust reddening could shift an AGN to the HII galaxy region in the diagram, since the $\beta$ would be observed to be lower, while $\beta$ would be similar."
 It is not. possible to estimate the reddening from the )almer emission lines. because the nuc2 s»ectrum is noisy.," It is not possible to estimate the reddening from the Balmer emission lines, because the $nuc2$ spectrum is noisy."
 However. the very low OLHL/L1L3—0.53 lis more suggestive of LIL galaxies than (tvpe 2 Αλ».," However, the very low $\beta$ $\pm$ 0.1 is more suggestive of HII galaxies than type 2 AGNs."
 Aloreover. its emission lines are very Harrow (FWILALS kms 1) compared with tvpical values of type 2 AGNs (several150 hundred kam ," Moreover, its emission lines are very narrow $\la$ 150 km $^{-1}$ ) compared with typical values of type 2 AGNs (several hundred km $^{-1}$ )."
3oth properties suggest that this is the actively star forming nucleus of a companion galaxy., Both properties suggest that this is the actively star forming nucleus of a companion galaxy.
 The shift in velocity relative to the quasar is )OcSNO km  (this error takes into account slit. ellect uncertainties. e See 82)," The shift in velocity relative to the quasar is $\pm$ 80 km $^{-1}$ (this error takes into account slit effect uncertainties, see $\delta$ 2)."
 This is within the range expected for mereing svstem (Atew hundred kin 1j , This is within the range expected for merging system $\la$ few hundred km $^{-1}$ ).
Both continuum and line emission are detected at the location of the knot and at the same z as the quasar., Both continuum and line emission are detected at the location of the knot and at the same $z$ as the quasar.
 The spectrum. is rather noisy and. 117 is not detected., The spectrum is rather noisy and $\beta$ is not detected.
 We can only sav that the location in the diagnostic clagrams (Lig., We can only say that the location in the diagnostic diagrams (Fig.
 3. bottom) using the LE? Dux upper limits overlaps with the LIL galaxies area and lies far from the standard AGN sequence.," 3, bottom) using the $\beta$ flux upper limits overlaps with the HII galaxies area and lies far from the standard AGN sequence."
 The knot looks sharper in the image than it does in, The knot looks sharper in the image than it does in
"the third. derivative of the lensing potential we have the unique combinations Fo= Fle = soo Μπ, - Jg"" ""Ezoe md. where the first flexion. J£. is a spin-l field and the new second Blexion. G. is seen to be a spin-3 field.","the third derivative of the lensing potential we have the unique combinations = | = ^* = = ^* = | = =, where the first flexion, ${\cal F}$, is a spin-1 field and the new second flexion, ${\cal G}$, is seen to be a spin-3 field."
 Here ὦ represents the position angle determining the direction of the vector or spin-3 component., Here $\phi$ represents the position angle determining the direction of the vector or spin-3 component.
 Expanding the Uexions in terms of the gradients of the shear field we find Fo- (ih. where the definition of the first Uexion agrees with our previous results in Section 2.," Expanding the flexions in terms of the gradients of the shear field we find = _1 _1 + _2 _2) + i _1 _2 = _1 _1 - _2 _2) + i _1 _2 _1), where the definition of the first flexion agrees with our previous results in Section 2."
 These two independent fields specify the weak “areiness” of the lensed image., These two independent fields specify the weak “arciness” of the lensed image.
 The complex representation allows us to [ind a consistency relation between the two Iexion fields. — which can be used as a check on measurements of F and C.," The complex representation allows us to find a consistency relation between the two flexion fields, ^* =, which can be used as a check on measurements of ${\cal F}$ and ${\cal G}$."
 We are also able to obtain a direct. description of the thirel order lensing tensor {δεν, We are also able to obtain a direct description of the third order lensing tensor $D_{ijk}$.
" Defining —J£,|i£» and C—O,|1G» we can then re-express ;,, as the sum of two terms Dijn=FindGijn. where the first (spin-1) term is and the second (spin-3) term is In order to obtain a visual understanding of the [exion quantities. we have used these forms for the jj), matrix in terms of F and G in order to calculate how a Gaussian image is transformed. by the various operations of weak lensing. according to equation (7))."," Defining $\flex = \flex_{1} + \mi \flex_2$ and $\sflex = \sflex_1 + \mi
\sflex_2$ we can then re-express $D_{ijk}$ as the sum of two terms $D_{ijk} = \flex_{ijk} + \sflex_{ijk}$, where the first (spin-1) term is and the second (spin-3) term is In order to obtain a visual understanding of the flexion quantities, we have used these forms for the $D_{ijk}$ matrix in terms of $\flex$ and $\sflex$ in order to calculate how a Gaussian image is transformed by the various operations of weak lensing, according to equation \ref{eq:2ndorder}) )."
 Phe results are shown in ligure 1. which displays the lensing operations in order of their spin properties.," The results are shown in Figure 1, which displays the lensing operations in order of their spin properties."
 The Gaussian galaxy is given a radius (standard deviation) of 1 arcsec: while the convergence and shear imposed on the galaxy are realistic in. each case). the Dexion is deliberately chosen to be extraordinarily large for visualisation purposes (0.28 aresec|. ef.," The Gaussian galaxy is given a radius (standard deviation) of 1 arcsec; while the convergence and shear imposed on the galaxy are realistic in each case), the flexion is deliberately chosen to be extraordinarily large for visualisation purposes (0.28 $^{-1}$, c.f."
 0.04 aresec+ intrinsic rms flexion on galaxies)., 0.04 $^{-1}$ intrinsic rms flexion on galaxies).
 We inimediately see the shapes induced by Hexion: the first Hlexion leads to a (vectorial. spin-1) skewness. while the second IHexion Leads to a threc-fold. (spin-3) shape.," We immediately see the shapes induced by flexion: the first flexion leads to a (vectorial, spin-1) skewness, while the second flexion leads to a three-fold (spin-3) shape."
 While the first flexion probes the local density via the eracdient of the shear field. the spin-3 second Uexion probes the nonlocal part of the gradient. of the shear field.," While the first flexion probes the local density via the gradient of the shear field, the spin-3 second flexion probes the nonlocal part of the gradient of the shear field."
 For example. consider a Schwarzschilel lens: the first Uexion is by definition zero everywhere except at the origin. as the eracient of the convergence is zero everywhere except at the origin.," For example, consider a Schwarzschild lens: the first flexion is by definition zero everywhere except at the origin, as the gradient of the convergence is zero everywhere except at the origin."
" However. there is certainly ""arciness"" generated by such a lens: this is described by the second exion."," However, there is certainly “arciness” generated by such a lens; this is described by the second flexion."
 We wil provide explicit expressions for the first ancl second IHexion generated by simple mass distributions in Sections 4 and 5., We will provide explicit expressions for the first and second flexion generated by simple mass distributions in Sections 4 and 5.
 The series of lensing distortions can clearly be continue to arbitrary. order by taking permutations of adcditiona spin-raising ancl lowering derivatives., The series of lensing distortions can clearly be continued to arbitrary order by taking permutations of additional spin-raising and lowering derivatives.
 For instance the nex order of distortion can be decomposed into three fields: a spin-4 field. GOVE. a spin-2 field. 0000v. and a spin- Geld. POI.," For instance the next order of distortion can be decomposed into three fields; a spin-4 field, $\de\de\de\de \psi$, a spin-2 field, $\de^* \de \de \de \psi$ , and a spin-0 field, $\de^*
\de^* \de\de \psi$ ."
 The n' order term can be decomposec into Int(1|»/2) independent spin fields with spins s=n. n 2.n 4d.-.O0 ibn is even or s+)1 if odd.," The $n^{th}$ order term can be decomposed into $(1+n/2)$ independent spin fields with spins $s=n$, $n-2$, $n-4$, $\cdots$, $0$ if $n$ is even or $\cdots 1$ if odd."
 Consistency relations similar to those for JF and G can be found for al 10 higher spin fields. which can also be used to estimate 10 convergence field. via Ixalser-Squires like relations (see Vlection 7).," Consistency relations similar to those for ${\cal F}$ and ${\cal G}$ can be found for all the higher spin fields, which can also be used to estimate the convergence field via Kaiser-Squires like relations (see Section 7)."
" ""FdDowever. in this paper we restrict ourselves to exploring 16 possibilities given by the first and second flexion."," However, in this paper we restrict ourselves to exploring the possibilities given by the first and second flexion."
 We will now proceed to calculate analytic expressions for both of the exion terms for simple lens models., We will now proceed to calculate analytic expressions for both of the flexion terms for simple lens models.
 In this section we present flexion predictions for galaxy-galaxy lensing uncer the assumption of a circularly svinmetric lens., In this section we present flexion predictions for galaxy-galaxy lensing under the assumption of a circularly symmetric lens.
 This is valid for a ealaxyv-ealaxy lensing approach where we do not reorient lens galaxies. resulting in à circularly averaged mean lens: in the following section we will consider the impact of having elliptical lenses.," This is valid for a galaxy-galaxy lensing approach where we do not reorient lens galaxies, resulting in a circularly averaged mean lens; in the following section we will consider the impact of having elliptical lenses."
 We consider a variety of dillerent. lens. models. and. show how flexion can be used to constrain them.," We consider a variety of different lens models, and show how flexion can be used to constrain them."
 The approximately Hat rotation curves observed in galaxies can be most simply reproduced by a model density profile which scales as pxr7., The approximately flat rotation curves observed in galaxies can be most simply reproduced by a model density profile which scales as $\rho \propto r^{-2}$.
 Such a profile can be obtained bv assuming a constant velocity dispersion for the clark matter throughout the halo. and so is known as the singular isothermal sphere (see c.g. Binney TremaineLOST).," Such a profile can be obtained by assuming a constant velocity dispersion for the dark matter throughout the halo, and so is known as the singular isothermal sphere (see e.g. Binney Tremaine1987)."
 The projected surface massdensity of the singular isothermal sphere (SIS) is, The projected surface massdensity of the singular isothermal sphere (SIS) is
"aby), helio-seinological data (D.B.Cueutherefαἱ )h and astercoscismological data from the pulsating white dwarf star GI117-B15À (ALBiesiadaefol) load to |GG=L10«10HeY for. <3.5 S.RayU.Mukhopadlivay. )).","), helio-seismological data \citeauthor{guenther}) ), and astereoseismological data from the pulsating white dwarf star G117-B15A \citeauthor{Biesiada}) ) lead to $\left|\dot{G}/G\right| \lessapprox 4.10 \times 10^{-11} yr^{-1}$, for $z\lesssim3.5$ \citeauthor{ray1}) )."
 Thus. in previous paper (M.R.Setareefab )). we investigated the holographic dark euergv scenario wider a varviug gravitational constant in the framework of IHorawa-Litshitz eravitv aud we extracted. the corresponding corrections to the dark energv equation-of-state Recently. a power-counting renormalizable. ultra- (UV) complete theory of gravity was proposed bv Totaava in LP.Uorava ..P.Horava ..P.Uorava )).," Thus, in previous paper \citeauthor{setare}) ), we investigated the holographic dark energy scenario under a varying gravitational constant in the framework of Horava-Lifshitz gravity and we extracted the corresponding corrections to the dark energy equation-of-state Recently, a power-counting renormalizable, ultra-violet (UV) complete theory of gravity was proposed by Hořaava in \citeauthor{hor2}, \citeauthor{hor1}, \citeauthor{hor4}) )."
" Althoueh preseutiug au infrared (IR) fixed point. namely General Relativity. in the UV the theory possesses a fixed point withau anisotropic. Lifshitz scaling between time and space of the form, αν»d4uw. toof where 6. 2.0! and f are the scaling factor. dynamical critical exponent. spatial coordinates and temporal coordinate. u the preseut work we are interested to study the Uolographic dark euergv in framework of Ioraava-Litshitz eravitv."," Although presenting an infrared (IR) fixed point, namely General Relativity, in the UV the theory possesses a fixed point withan anisotropic, Lifshitz scaling between time and space of the form $x^{i}\to\ell~x^{i}$, $t\to\ell^z~t$, where $\ell$, $z$, $x^{i}$ and $t$ are the scaling factor, dynamical critical exponent, spatial coordinates and temporal coordinate, n the present work we are interested to study the Holographic dark energy in framework of Ho\v{r}aava-Lifshitz gravity."
 We extend previous study (M.R.Setareetal )) to the non-flat case and we want consider 1e effectof curvature constant of uou flat space on the resalts which is obtained im the flat space. it is secu that ie resonable range of A Eq.(3)) iu nou flat space is arecr than auswer rauge of À in flat space.," We extend previous study \citeauthor{setare}) ) to the non-flat case and we want consider the effectof curvature constant of non flat space on the resalts which is obtained in the flat space, it is seen that the resonable range of $\lambda$ \ref{2g}) ) in non flat space is larger than answer range of $\lambda$ in flat space."
 The paper is organized as follows: in the next section we found 1¢ holographic dark cucrey with eravitation coustant aepeud on time and derive the differeutial equation iat specify the evolution of dark energy parameter., The paper is organized as follows: in the next section we found the holographic dark energy with gravitation constant depend on time and derive the differential equation that specify the evolution of dark energy parameter.
 Iu Sec., In Sec.
 3. we obtain the parameter of the dark energy equation of state at the low redshift.," 3, we obtain the parameter of the dark energy equation of state at the low redshift."
Eveutuallv. the atter section is devoted to conclusion.,"Eventually, the latter section is devoted to conclusion."
 Tu the case where the space-time ecometry is a non-fat Robertson-Walkser: with e(f) the scale factor. in comoving coordinates Gor0.5). where & denotes the spacial curvature with hos1.0.1 correspondiug to open. flat and closed universe respectively.," In the case where the space-time geometry is a non-flat Robertson-Walker: with $a(t)$ the scale factor, in comoving coordinates $(t,r,\theta,\varphi)$, where $k$ denotes the spacial curvature with $k=-1,0,1$ corresponding to open, flat and closed universe respectively."
" In this case. the first Priedinaun equation im the framework of IHoraava-Lifslütz eravity writes (ALR.Setaree£al): where. ϱPin|pa is the energy density. py,= pa are dark matter deusity and dark cucrey density respectively, A is a dineusioual constant and j—NON.sinLIF."," In this case, the first Friedmann equation in the framework of Hořaava-Lifshitz gravity writes \citeauthor{setare}) ): where, $\rho=\rho_m+\rho_{\Lambda}$ is the energy density, $\rho_m=\rho_{m_0} a^{-3}$ and $\rho_{\Lambda}$ are dark matter density and dark energy density respectively, $\lambda$ is a dimensional constant and $\beta=\frac{\kappa^4\mu^2\Lambda }{8(3\lambda-1)^2}$."
" Mere απD=οπέςdE and p,,, indicate.- the present value that quantitv."," Here $\kappa^2=8\pi G$, and $\rho_{m_0}$ indicate the present value that quantity."
 Then. from the Eq.(3)). we introduce the effective density parameter ον— substituting:4iIos ," Then, from the \ref{2g}) ), we introduce the effective density parameter $\Omega_{\Lambda_e}=\frac{\Omega_{\Lambda}}{2\left( 3\lambda-1\right)}\equiv \frac{8\pi G}{6\left( 3\lambda-1\right) }\frac{\rho_{\Lambda}}{H^2}$ ."
theMin Eq.(1))1 into: OO4 we obtain: where ος=¢VE(BA 1)., Substituting the \ref{1e}) ) into $\Omega_{\Lambda}$ we obtain: where $c_e=c/\sqrt{2\left(3\lambda-1 \right)} $ .
 Iu this case. the cosmological leneth L in (1)) is considered to be (OQ.C. where and dy is the future event lorizou defiud by Ueuceforth. we will use lue. as an independent variable.," In this case, the cosmological length $L$ in \ref{3g}) ) is considered to be \citeauthor{nonflat}) ): where and $d_{h}$ is the future event horizon defind by Henceforth, we will use $\ln a$, as an independent variable."
" Therefore. we define. .X=i: wand Y. Ja. κο that X= XII. A straightforward calculation using (0) and (3)) leads to (AL.Jamul 3): where Iu order to clarity the effect of variety Con the O4, we should to eet rid of 11 into Eq. (8))."," Therefore, we define, $\dot{X}=\frac{dX}{dt}$ , and $X'=\frac{dX}{d\ln a}$ , so that $\dot{X}=X'H$ A straightforward calculation using \ref{3g}) ) and \ref{2g}) ) leads to \citeauthor{Jamil:2009sq}) ): where In order to clarify the effect of variety Gon the $\Omega_{\Lambda e}$, we should to get rid of $\dot{H}$ into Eq. \ref{5g}) )."
 Iu this regard. differentiation of the Friedinaun equation eive rise to," In this regard, differentiation of the Friedmann equation give rise to"
related temperature profiles.,related temperature profiles.
" We choose a logarithmic x-axis, thus focusing on the center of the simulation box."," We choose a logarithmic x-axis, thus focusing on the center of the simulation box."
" As an immediate effect of the heating due to the photoionizing UV background, the configuration is heated up to temperatures of T«2x10* K during the reionization at redshift z~6."," As an immediate effect of the heating due to the photoionizing UV background, the configuration is heated up to temperatures of $T \approx 2 \times 10^4$ K during the reionization at redshift $z\approx6$."
 This results in a pressure several orders of magnitudes higher than in the non-radiative simulations., This results in a pressure several orders of magnitudes higher than in the non-radiative simulations.
" Therefore, the adiabatic collapse before redshift (z& 1) does not produce one single peak, but an supported by pressure."," Therefore, the adiabatic collapse before redshift $z\approx1$ ) does not produce one single peak, but an supported by pressure."
 The further evolution now depends on the size of the perturbation length scale L., The further evolution now depends on the size of the perturbation length scale $L$.
" For the smallest perturbation scale L=2 Mpc the speed of sound inside this core remains always higher than the infall velocity, and therefore the shock cannot form anymore."," For the smallest perturbation scale $L=2$ Mpc the speed of sound inside this core remains always higher than the infall velocity, and therefore the shock cannot form anymore."
" The whole profile, now sustained by the pressure of the gas only, is more extended than in the non-radiative case."," The whole profile, now sustained by the pressure of the gas only, is more extended than in the non-radiative case."
 For L>2 Mpc the infall velocity becomes higher than the sound velocity at some moment (this can be obtained from the scaling considerations discussed above)., For $L>2$ Mpc the infall velocity becomes higher than the sound velocity at some moment (this can be obtained from the scaling considerations discussed above).
" Thus, a shock is able to form."," Thus, a shock is able to form."
" This shock is not generated in the center, but forms at the edges of the pre-shock core."," This shock is not generated in the center, but forms at the edges of the pre-shock core."
" From there, it moves outward, like in the non-radiative case."," From there, it moves outward, like in the non-radiative case."
" Additionally, a fan-like wave penetrates into the core, effectively shrinking its size."," Additionally, a fan-like wave penetrates into the core, effectively shrinking its size."
" The whole configuration can be separated into an inner isothermal core, a shocked region of higher temperatures, and an outer region at low density and low temperature."," The whole configuration can be separated into an inner isothermal core, a shocked region of higher temperatures, and an outer region at low density and low temperature."
 The size of the core is decreasing with increasing length scale L., The size of the core is decreasing with increasing length scale $L$.
" Outside of the core region, the results of the simulations differ only slightly from the non-radiative case."," Outside of the core region, the results of the simulations differ only slightly from the non-radiative case."
 Especially the position of the outer shock appears to be unaffected., Especially the position of the outer shock appears to be unaffected.
" A special situation occurs in the L=4 Mpc simulation, where an effective outflow can be noticed, which appears as a positive density flux the core."," A special situation occurs in the $L = 4$ Mpc simulation, where an effective outflow can be noticed, which appears as a positive density flux the core."
" This is caused by the lower density inside the core compared to the runs with higher L, resulting in longer cooling times, and thus a less effective dissipation of the energy input by the further infall."," This is caused by the lower density inside the core compared to the runs with higher $L$, resulting in longer cooling times, and thus a less effective dissipation of the energy input by the further infall."
 The influence of the perturbation scale on the propertiesof the core will be further examined in Sect. ??.., The influence of the perturbation scale on the propertiesof the core will be further examined in Sect. \ref{sScaling}. .
photometric calibration was done bv comparing with the corresponding SDSS images in the r filter.,photometric calibration was done by comparing with the corresponding SDSS images in the $r$ filter.
 The l-o brightness [Iuctuation of the resulting images is 27.0-27.5 mag 7., The $\sigma$ brightness fluctuation of the resulting images is 27.0-27.5 mag $^{-2}$.
 The spectroscopic observations were taken αἱ the 104m GTC (ORAL La Palma. Spain).," The spectroscopic observations were taken at the 10.4 m GTC (ORM, La Palma, Spain)."
 Details of the telescope and the instrumental configuration can be found on the telescope web site (vwww.gtc.iac.es)., Details of the telescope and the instrumental configuration can be found on the telescope web site (www.gtc.iac.es).
 We used (he OSIRIS spectrograph with the R1000DB erism., We used the OSIRIS spectrograph with the R1000B grism.
 The observations were carried out in August 2010 in service mode., The observations were carried out in August 2010 in service mode.
 A single position ol a long slit crossing the two members of each CG candidate with unknown recshilt was taken., A single position of a long slit crossing the two members of each CG candidate with unknown redshift was taken.
 The spatial sampling was 0.25 arcsec/pixel., The spatial sampling was 0.25 arcsec/pixel.
 The slit width was 1 arcsec., The slit width was 1 arcsec.
 The nights were clear with (races of dust ancl light cirrus., The nights were clear with traces of dust and light cirrus.
 For wavelength calibration. we used Igàár and Ne lamps taken at the end of each night.," For wavelength calibration, we used HgAr and Ne lamps taken at the end of each night."
 The stabilitv of the waveleneth calibration during the night was checked with the main skv lines., The stability of the wavelength calibration during the night was checked with the main sky lines.
 The effective spectral resolution was 7Α, The effective spectral resolution was 7.
"ι, The spectra were analyzed following a standard procedure using IRAF. which comprises bias sublraction. flat [ield correction. coaddilion of exposures of the same field. wavelength: calibration. and extraction of the spectra."," The spectra were analyzed following a standard procedure using IRAF, which comprises bias subtraction, flat field correction, coaddition of exposures of the same field, wavelength calibration, and extraction of the spectra."
 The flat field correction was clone onlv in the red part of each spectrum due to the low response of the calibration lamp in the blue part (this correction is ~ 1/200).," The flat field correction was done only in the red part of each spectrum due to the low response of the calibration lamp in the blue part (this correction is $\sim
1/200$ )."
 The spectral calibration provides a sampling of 2.07 and is quite aceurate (~0.05 Aj). as checked through the position of the atmospheric OI line.," The spectral calibration provides a sampling of 2.07 and is quite accurate $\sim 0.05$ ), as checked through the position of the atmospheric OI line."
 We used standard spectroscopic stars [rom the catalog bv Oke (1990) to correct for the response of the configuration at different wavelengths., We used standard spectroscopic stars from the catalog by Oke (1990) to correct for the response of the configuration at different wavelengths.
 This does not provide an absolute flux calibration owing to the conditions of the observations. the different orientation of the slit width. ete.," This does not provide an absolute flux calibration owing to the conditions of the observations, the different orientation of the slit width, etc."
 We did not correct for the tellurie band at 7600., We did not correct for the telluric band at 7600.
À.. Figure l shows images of each CG candidate aud (he extracted spectra of (heir menibers (for completeness. we also include the spectra of the galaxies observed with," Figure 1 shows images of each CG candidate and the extracted spectra of their members (for completeness, we also include the spectra of the galaxies observed with"
Correlation functions are some of the most commonly used statistics in cosmology.,Correlation functions are some of the most commonly used statistics in cosmology.
 They have a long history in quantifying the clustering of galaxies in the Universe (see Peebles 1980), They have a long history in quantifying the clustering of galaxies in the Universe (see Peebles 1980).
 There is a hierarchy of correlation functions., There is a hierarchy of correlation functions.
 The two-point correlation function (2PCF) compares the number of pairs of data points. as a function of separation. with that expected from a Poisson distribution.," The two–point correlation function (2PCF) compares the number of pairs of data points, as a function of separation, with that expected from a Poisson distribution."
 Next in the hierarchy is the 3—point correlation function (3PCF). which compares the number of data triplets. as a function of their triangular-configuration. to that expected from Poisson.," Next in the hierarchy is the 3–point correlation function (3PCF), which compares the number of data triplets, as a function of their triangular–configuration, to that expected from Poisson."
 Higher-order correlations are detined analogously., Higher-order correlations are defined analogously.
 As discussed by many authors. the higher-order correlation functions contain a variety of important cosmological information. which complements that from the 2PCF (Groth.&Peebles1977:Balian&Schaeffer 1989).," As discussed by many authors, the higher–order correlation functions contain a variety of important cosmological information, which complements that from the 2PCF \citep{GP1977,BS1989}."
. These include tests of Gaussianity and the determination of galaxy bias as a function of scale (Suto1993:Kayoetal.2004:Lahav&Suto 2004).," These include tests of Gaussianity and the determination of galaxy bias as a function of scale \citep{Suto1993,JB1998,TJ2003,JB2004,Kayo2004,LS2004}."
". Such tests can also be performed using the Fourier-space equivalent of the vun the bi-spectrum (Peebles1980:Scoccimarroal.1999,2001:Verdeet2002) or other statistics such as the etvoid ο distribution and Minkowski functionals(Meckeetal.1994)."," Such tests can also be performed using the Fourier-space equivalent of the 3PCF, the bi-spectrum \citep{PEEBLES1980,Sc1999,PSCz,Verde2002} or other statistics such as the void probability distribution and Minkowski \citep{MBW1994}."
. Recent results from these complementary statistics using the SDSS main galaxy sample include Hikageetal.(2002.2003.2005) and (2005)," Recent results from these complementary statistics using the SDSS main galaxy sample include \cite{H2002,H2003,H2005} and \cite{Park2005}."
" While the 3PCF is easier to correct for survey edge effects than these other statistics. measurements of the 3PCF have been limited by the availability of large redshift surveys of galaxies (see Szapudi. Meiksin Nichol 1996. Frieman Gaztanaga 1999, Szapudi et al."," While the 3PCF is easier to correct for survey edge effects than these other statistics, measurements of the 3PCF have been limited by the availability of large redshift surveys of galaxies (see Szapudi, Meiksin Nichol 1996, Frieman Gaztanaga 1999, Szapudi et al."
 2002 for 3PCF analyses of large solid angle catalogues of galaxies) and the potentially prohibitive computational time needed to count all possible triplets of galaxies (naively. this count scales as O(N?). where NV is the number of galaxies in the sample).," 2002 for 3PCF analyses of large solid angle catalogues of galaxies) and the potentially prohibitive computational time needed to count all possible triplets of galaxies (naively, this count scales as $O(N^3)$, where $N$ is the number of galaxies in the sample)."
 In this paper. we resolve these two problems through the application a new N-point correlation function algorithm (Mooreetal.ofσος.2001) to the galaxy data of the Sloan Digital Sky Survey York et al.," In this paper, we resolve these two problems through the application of a new N–point correlation function algorithm \citep{MOORE2001} to the galaxy data of the Sloan Digital Sky Survey (SDSS; York et al."
 2000)., 2000).
 We present herein measurements of the 3PCF from the SDSS main galaxy sample., We present herein measurements of the 3PCF from the SDSS main galaxy sample.
 Our measurements illustrate the sensitivity of the 3PCF to known large-scale structures in the SDSS (Gottetal.2005)., Our measurements illustrate the sensitivity of the 3PCF to known large-scale structures in the SDSS \citep{Gott2005}.
.. They are complementary to the work of Kayoetal.(2004) who explicitly explored the luminosity and morphological dependence9 of the 3PCF using SDSS volume—limited galaxy samples., They are complementary to the work of \cite{Kayo2004} who explicitly explored the luminosity and morphological dependence of the 3PCF using SDSS volume–limited galaxy samples.
 These measurements of the 3PCF will help facilitate constraints on the biasing of galaxies and will aid in the development of theoretical predictions for the higher-order functions (Scoccimarroetal.2001:Takada&J," These measurements of the 3PCF will help facilitate constraints on the biasing of galaxies and will aid in the development of theoretical predictions for the higher–order correlation functions \citep{Sc2001,TJ2003}."
ain200 correlation paper. we use the dimensionless Hubble constant S.=LodM‘heOkms+MpcLH the matter density parameter {δι=0.3. dimensionless cosmological constant 4.=0.7. unless stated otherwise.," Throughout this paper, we use the dimensionless Hubble constant $h \equiv H_{\rm 0}/100\,{\rm km\,s^{-1}\,Mpc^{-1}}$, the matter density parameter $\Omega_{\rm m}=0.3$, and the dimensionless cosmological constant $\Omega_\Lambda=0.7$, unless stated otherwise."
 To facilitate the rapid ealeulation of the higher-order correlation functions. we have designed and implemented a new N-point correlation function (NPCF) algorithm based on k-d trees. which are multi-dimensional binary search tree for points in a dimensional space.," To facilitate the rapid calculation of the higher–order correlation functions, we have designed and implemented a new N–point correlation function (NPCF) algorithm based on k-d trees, which are multi–dimensional binary search tree for points in a k-dimensional space."
 The k-d tree is composed of a series of inter— nodes. which are created by recursively splitting each node along its longest dimension. thus creating two smaller child nodes.," The k-d tree is composed of a series of inter--connected nodes, which are created by recursively splitting each node along its longest dimension, thus creating two smaller child nodes."
 This recursive splitting is stopped when a pre-determined number of data points is reached in each node (we used <=20 data points herein}., This recursive splitting is stopped when a pre-determined number of data points is reached in each node (we used $\leq20$ data points herein).
 For our NPCF algorithm. we used an enhanced version of the k-d tree technology. namely multi-resolutional k-d trees with cached statistics (mrkdtree). which store additional statistical information about the search tree. and the data points in each node.¢.g.. we store the total count and centroid of all data in each node.," For our NPCF algorithm, we used an enhanced version of the k-d tree technology, namely multi–resolutional k-d trees with cached statistics (mrkdtree), which store additional statistical information about the search tree, and the data points in each node, we store the total count and centroid of all data in each node."
 The key to our NPCF algorithm is to use multiple mrkdtrees together. and store them in main memory of the computer (rather than on disk). to represent the required N-point function.¢.¢.. we use 3 mrkdtrees to compute the 3PCF. 4 mrkdtrees for the 4PCF. and so on.," The key to our NPCF algorithm is to use multiple mrkdtrees together, and store them in main memory of the computer (rather than on disk), to represent the required N–point function, we use 3 mrkdtrees to compute the 3PCF, 4 mrkdtrees for the 4PCF, and so on."
 The computational efficiency is increased by pruning these trees wherever possible. and by using the cached statistics on the tree as much as possible.," The computational efficiency is increased by pruning these trees wherever possible, and by using the cached statistics on the tree as much as possible."
 The details of mrkdtrees and our NPCF algorithm (known as i have already s outlined in several papers (Mooreetal.2001:Nichol2003:Grayal.2004).," The details of mrkdtrees and our NPCF algorithm (known as ) have already been outlined in several papers \citep{MOORE2001,NICHOL2003,GRAY2004}."
. Similar tree-based πηραον now ave been discussed by Szapudietal.(2001}., Similar tree–based computational algorithms have been discussed by \cite{Szapudi2001}.
. The details of the SDSS survey are given in a series of technical papers M Fukugitaetal.neGunnnYork Blantonκ“.," The details of the SDSS survey are given in a series of technical papers by \cite{fuk96, gunn98, Y2000, hogg01, Strauss2002, smith02,
 pier03, B2003b, ivezic04, DR3}. ."
For the computations discussed herein. we use two SDSS ," For the computations discussed herein, we use two SDSS catalogues."
"The first is a volume-limited sample of 36738 galaxies in the redshift range of 0.05xz0.095 and absolute magnitude range of 3239Musis20.5 (for fh=0.7 and the >=0.0 SDSS + filter. or """"rin Blantonetal.(2003b) 3). covering 2364 deg of the SDSS photometric survey."," The first is a volume–limited sample of 36738 galaxies in the redshift range of $0.05 \le z \le
0.095$ and absolute magnitude range of $-23 \le {M_{^{0.0}r}} \le
-20.5$ (for $h=0.7$ and the $z=0.0$ SDSS $r$ filter, or $^{0.0}r$ in \cite{B2003b} ), covering 2364 $^2$ of the SDSS photometric survey."
 All the magnitudes were reddening corrected using Schlegel.Finkbbeiner. (1998). and the software (Blanton 2003b).," All the magnitudes were reddening corrected using \cite{SFD98}, and the software \citep{B2003b}."
. The second “sample is the same as “Sample 127 used by Popeetal.(2004) and contains 134741. galaxies over 2406 deg., The second sample is the same as “Sample 12” used by \cite{Pope2004} and contains 134741 galaxies over 2406 $^2$.
" This latter sample is not volume-limited. but is constrained to the absolute magnitude range of |22<Moi,;19 tor APx15 for h= 1. M» MEE the.=""M Ὃ r filter system. ortr magnitudes)(Blantonetvolumeal.20030:Zehavi05)."," This latter sample is not volume–limited, but is constrained to the absolute magnitude range of $-22 \le {M_{^{0.1}r}}
\le -19$ (or $M^{*}\pm1.5$ magnitudes) for $h=1$ , and using the $z=0.1$ SDSS r filter system, or $^{0.1}r$ \citep{B2003b,Zehavi2005}."
" To compare the two samples. our itited sample "" the absolute magnitude range of =23.54<Moi,21.04 in the same ""ty filter as used for the Pope et al."," To compare the two samples, our volume–limited sample has the absolute magnitude range of $-23.54 \le {M_{^{0.1}r}} \le -21.04$ in the same $^{0.1}r$ filter as used for the Pope et al."
" sample: assuming a conversion of ic""|0.23 for the SDSS main galaxy sample with a median color at >=0.0 of """"(g—r)&OLS."," sample; assuming a conversion of ${^{0.1}}r\simeq {^{0.0}}r + 0.23$ for the SDSS main galaxy sample with a median color at $z=0.0$ of ${^{0.0}}(g-r)\simeq
0.8$."
 This gives a mean space of density5.25.10.?5? * which is comparable to the space densities of the SDSS main galaxy sample given in Table 2 of Zehavietal.(2005," This gives a mean space density of $8.25\times 10^{-3}\,h^{3}$ $^{-3}$, which is comparable to the space densities of the SDSS main galaxy sample given in Table 2 of \cite{Zehavi2005}."
 We have made no correction for missing galaxies due to fibre (i.e.. two SDSS fibres can notbe placed closer than 55 areseconds on the sky).," We have made no correction for missing galaxies due to fibre--collisions (i.e., two SDSS fibres can notbe placed closer than 55 arcseconds on the sky)."
 We do not expectthis observational constraint to bias our correlation functions as the adaptive tilting of SDSS spectroscopic plates reduces the problem to τά of possible target galaxies being missed (see Blanton et al., We do not expectthis observational constraint to bias our correlation functions as the adaptive tilting of SDSS spectroscopic plates reduces the problem to $\simeq7\%$ of possible target galaxies being missed (see Blanton et al.
 2003a for details)., 2003a for details).
 Furthermore. this bias will only affect pairs of galaxies," Furthermore, this bias will only affect pairs of galaxies"
doninates over the iucrease in density because of the high ionization poteutial of IHTe-Ilike Si NTT.,dominates over the increase in density because of the high ionization potential of He-like Si XIII.
 As secu in Figure 9 similar treuds are seeu in oxyeen., As seen in Figure \ref{fig:charge} similar trends are seen in oxygen.
 Tere. however. even a the highest efBcieucies iu the high density cases. oxyeeu is alanost fully iouizec.," Here, however, even at the highest efficiencies in the high density cases, oxygen is almost fully ionized."
 Equating these results to Figures 7 auc ὃ mcaus tha in the low clensity case (Fig. 7)).," Equating these results to Figures \ref{fig:te_p1cc} and \ref{fig:te_1cc} means that in the low density case (Fig. \ref{fig:te_p1cc}) ),"
 the lower density Προς ess collisional ionization aud less cfiicicut Coulomb jieafing. so that the plastua stavs less jonized.," the lower density implies less collisional ionization and less efficient Coulomb heating, so that the plasma stays less ionized."
 More xeciselv. iu the low deusitv models. the Ie-like states are not reached. aud forbidden line euission is domunatcc winner shell excitation of lower ious.," More precisely, in the low density models, the He-like states are not reached, and forbidden line emission is dominated by inner shell excitation of lower ions."
 In higher deusity uodels. the IHe-like states are reached. and the expectec dependence with temperature is seen (c.f.," In higher density models, the He-like states are reached, and the expected dependence with temperature is seen (c.f."
 Fig. 8))., Fig. \ref{fig:te_1cc}) ).
 We stress that in the low deusitv models. some elements. such as silicon. will be highly uuderiouized aud the II-ike triplets will be quite faint.," We stress that in the low density models, some elements, such as silicon, will be highly underionized and the He-like triplets will be quite faint."
" These lines will be lax o detect, and when combined with any nonuthenua contimmiun cussion. they might be undetectable with current mstruiueutatiou."," These lines will be hard to detect, and when combined with any nonthermal continuum emission, they might be undetectable with current instrumentation."
 We cud this section with Figure 10.. a plot of the variation im GY. average charge state. aud electron teiiperature as a function of ambicut imiediun doeusity. at coustaut acceleration efficiencies.," We end this section with Figure \ref{fig:consteff}, a plot of the variation in $G'$, average charge state, and electron temperature as a function of ambient medium density, at constant acceleration efficiencies."
 This plot shows how these values vary between the test particle case and a ecffiiient shock., This plot shows how these values vary between the test particle case and a efficient shock.
" As expected. the average electron teiiperature in the test particle case is higher than iu the efficient case, but the average charee states are very similar."," As expected, the average electron temperature in the test particle case is higher than in the efficient case, but the average charge states are very similar."
" The biggest differences come frou the ratio of C"". where in the low deusitv models. the G ratio is Iigher in the test particle case than in the cfiicient case. while above a density of < 1.0 7. the relation is reversed. with the C ratio in the efficient. case is higher."," The biggest differences come from the ratio of $G'$, where in the low density models, the $G'$ –ratio is higher in the test particle case than in the efficient case, while above a density of $\la$ 1.0 $^{-3}$, the relation is reversed, with the $G'$ –ratio in the efficient case is higher."
 Frou. this plot. it is clear that at identical ambient medium densities. simular charge states can result between the efficicut andiuefficient models. while the tempcraturcs can differ bv ~ and the line ratios can differ bv ~105..," From this plot, it is clear that at identical ambient medium densities, similar charge states can result between the efficient andinefficient models, while the temperatures can differ by $\sim$ and the line ratios can differ by $\sim$."
 For instance. for au ambicut mediuui density of 1.0 +. the weighted temperature in the efficient model is lower than in the test particle case.," For instance, for an ambient medium density of 1.0 $^{-1}$, the weighted temperature in the efficient model is lower than in the test particle case."
 However. the ionization state for silicon here is virtually ideutical. but the G ratio in the efficient model is higher than in the inefficient model.," However, the ionization state for silicon here is virtually identical, but the $G'$ –ratio in the efficient model is higher than in the inefficient model."
 At tempcraturcs which can differ bv as much as they do in Figure 10.. what we see here is that at constant ambicut nmediuuu density but differing efiicieucies. different lines are seen at different inteusites.," At temperatures which can differ by as much as they do in Figure \ref{fig:consteff}, what we see here is that at constant ambient medium density but differing efficiencies, different lines are seen at different intensites."
 Since the line ratios aud underline electron temperatures are different between the efücieut and incficicnt models. the shape of the resultant spectra will differ.," Since the line ratios and underlying electron temperatures are different between the efficient and inefficient models, the shape of the resultant spectra will differ."
 The models preseuted here differ from existing available models that compute the A-ray emission from shocks du several wavs. but the relevant cdiffereuce here is that the nonequilibrunm ionization is followed siauultaneouslv with the shock cinamics and particle acceleration.," The models presented here differ from existing available models that compute the X-ray emission from shocks in several ways, but the relevant difference here is that the nonequilibrium ionization is followed simultaneously with the shock dynamics and particle acceleration."
 The ionization vector is then passed directly to an cuussivity code to compute the thermal N-rav Cluission., The ionization vector is then passed directly to an emissivity code to compute the thermal X-ray emission.
 Since the NET is evolved simultancously with the lvdvodvueamic evolution. our calculation includes changes to the NET as a result of adiabatic losses. or as is the case in an efficient model. increased conrpressiou belind the shock resultiug from efücieut shock acceleration.," Since the NEI is evolved simultaneously with the hydrodynamic evolution, our calculation includes changes to the NEI as a result of adiabatic losses, or as is the case in an efficient model, increased compression behind the shock resulting from efficient shock acceleration."
 To illustrate this effect. in Figure 11.. we compare the thermal ταν ecuuission from models where we follow the time-dependent ionization with the lydrodvuamics (black curves} to the thermal X-ray cussion from models where we calculate the ionization structure from the final ionization age (7) aud electron temperature (red curves).," To illustrate this effect, in Figure \ref{fig:comp}, we compare the thermal X-ray emission from models where we follow the time-dependent ionization with the hydrodynamics (black curves) to the thermal X-ray emission from models where we calculate the ionization structure from the final ionization age $\tau$ ) and electron temperature (red curves)."
 Here. the final ionization for each cell age is defined as the 7 at the end of the simulation.," Here, the final ionization for each cell age is defined as the $\tau$ at the end of the simulation."
 Then. for conrparison. we reun the NET calculation for the final electron temperature and the time equivalent to the age of the SNR at the cud of the lvdrodvuamical simulation.," Then, for comparison, we re-run the NEI calculation for the final electron temperature and the time equivalent to the age of the SNR at the end of the hydrodynamical simulation."
 For these mnodels we do not include amy efficient particle acceleration. as that will ouly complicate the interpretation of these results.," For these models, we do not include any efficient particle acceleration, as that will only complicate the interpretation of these results."
" Computing the ionization structure in each exid cell from the final 7 is similar to assuniue a suele 7. sinele T; model for cach cell In (e.g. the model). except that in our conrputation of 7 and T, in the lywdrocdwuamics. we still account for adiabatic cooling iu the shocked eas. since these values were computed initially from the lydrodvuamiics."," Computing the ionization structure in each grid cell from the final $\tau$ is similar to assuming a single $\tau$, single $T_e$ model for each cell in (e.g. the model), except that in our computation of $\tau$ and $T_e$ in the hydrodynamics, we still account for adiabatic cooling in the shocked gas, since these values were computed initially from the hydrodynamics."
 This approach is similar to Ellisonctal.(2007)., This approach is similar to \citet{ellison07}.
.. As secu iu Figure 1l. there are differences between the two resultaut spectra.," As seen in Figure \ref{fig:comp}, there are differences between the two resultant spectra."
 In the low deusitv (yg = Ul cn?) model the caleulation of the cluitted spectrum from the final ionization age sliebtlv overpredicts the final thermal X-ray cussion. lost evidently seen around the Ie-like lines of maguesiuu and silicon. as well as from the emission of L-shell iron.," In the low density $_{p,0}$ = 0.1 $^{-3}$ ) model, the calculation of the emitted spectrum from the final ionization age slightly overpredicts the final thermal X-ray emission, most evidently seen around the He-like lines of magnesium and silicon, as well as from the emission of L-shell iron."
 Not surprisingly. the shape of the uuderling continua are consistent.," Not surprisingly, the shape of the underlying continua are consistent."
 This is expected because the shape of the continmuu is determined by the electron temperature. which are identical between the two models.," This is expected because the shape of the continuum is determined by the electron temperature, which are identical between the two models."
" Iu the higher density (a,.u = L0 cmj model. the differences are mmch larger. particularly at low enerev."," In the higher density $_{p,0}$ = 1.0 $^{-3}$ ) model, the differences are much larger, particularly at low energy."
 Iu particular. the ionization state aud thermal N-ray cinission computed from the final ionization age and temperature muderprecdict the Fe-L cunission.," In particular, the ionization state and thermal X-ray emission computed from the final ionization age and temperature underpredict the Fe-L emission."
 We interpret this cditferenuce as beime a direct result of uot following the ionization explicitly: that is. when colmputing the ionization state from a single density and Af. intermediary ionization states are not correctly populated. since the calculation asstues that the shocked gas has gone from a cold (104 K) eas to hot (several 109 IS) in a suele timestep.," We interpret this difference as being a direct result of not following the ionization explicitly; that is, when computing the ionization state from a single density and $\Delta t$, intermediary ionization states are not correctly populated, since the calculation assumes that the shocked gas has gone from a cold $^{4}$ K) gas to hot (several $^{6-7}$ K) in a single timestep."
 We note that were a spectruni Hike this ft with existing models. it would appear that the shocked. plasma is overabundaut in metals such as argon and calcium.," We note that were a spectrum like this fit with existing models, it would appear that the shocked plasma is overabundant in metals such as argon and calcium."
 Figure 11 clearly demonstrates the need to compute the ionization sclf&consistently with the shock lvdrodwuamics (regardless of whether one considers the effects of efficieut. diffusive shock acceleration)., Figure \ref{fig:comp} clearly demonstrates the need to compute the ionization self-consistently with the shock hydrodynamics (regardless of whether one considers the effects of efficient diffusive shock acceleration).
 Iu Figure 13.. we plot the observed ditfereuces between the thermal enüssioun from a test-particle model aud the combined thermal and nouthermal cussion from an efficient model.," In Figure \ref{fig:tp_vs_eff}, we plot the observed differences between the thermal emission from a test-particle model and the combined thermal and nonthermal emission from an efficient model."
" The black crosses in Figure 19 are sSiuaulated data for a 50 ksec Chandra ACTS-S observation of au SNR with eps, = 0 aud νο = 0.2 ."," The black crosses in Figure \ref{fig:tp_vs_eff} are simulated data for a 50 ksec Chandra ACIS-S observation of an SNR with $\effDSA$ = 0 and $n_{p,0}$ = 0.3 $^{-3}$ ."
 The blue curve is the nouthermal emission. the red curve is the thermal cussion and the black curve is their stun.," The blue curve is the nonthermal emission, the red curve is the thermal emission and the black curve is their sum."
" In the efficient model we assunie eps, = aud yy = 0.3 Ὁ, "," In the efficient model, we assume $\effDSA$ = and $n_{p,0}$ = 0.3 $^{-3}$ ."
As expected. at the high energy," As expected, at the high energy"
ab d.rjan.,at $4.7\micron$.
 For simplicity. no uncerlving contin was assumed in fitting the line emission.," For simplicity, no underlying continuum was assumed in fitting the line emission."
 This is reasonable exgiven our limited egoal of estimatinge the radial extent of the emittinge 0gas., This is reasonable given our limited goal of estimating the radial extent of the emitting gas.
 An important model parameter is the stellar mass., An important model parameter is the stellar mass.
 We can estimate the stellar mass of V836 Tau both from pre-main-sequence evolutionary (racks and from available dynamical nass estimates of pre-iain-sequence stars of the same spectral (vpe and similar age., We can estimate the stellar mass of V836 Tau both from pre-main-sequence evolutionary tracks and from available dynamical mass estimates of pre-main-sequence stars of the same spectral type and similar age.
 The Ki 1 spectral (wpe (or Teg=4000zx200 IxXIx: White IHillenbrand 2004) corresponds to a mass of 0.71uM. using the evolutionary tracks of Siess et ((2000)., The K7 $\pm 1$ spectral type (or $_{\rm eff} = 4000\pm 200$ K; White Hillenbrand 2004) corresponds to a mass of $0.71^{+0.24}_{-0.18} \Msun$ using the evolutionary tracks of Siess et (2000).
 Similarly. other T Taur stus in Taurus with AT spectral (vpes have measured dvnanmical masses (Simon et 22000) of 0.72.4. (DL Tan) and 0.8437. (GM Aur).," Similarly, other T Tauri stars in Taurus with K7 spectral types have measured dynamical masses (Simon et 2000) of $0.72\Msun$ (DL Tau) and $0.84\Msun$ (GM Aur)."
 Thus. the ντ dynamical masses and IIR. diagram position of V836 Tau are consistent. with a stellar mass of 0.7—0.511...," Thus, the K7 dynamical masses and HR diagram position of V836 Tau are consistent with a stellar mass of $0.7-0.8 \Msun$."
 If the spectral type is larger or smaller by one subclass. a larger range in mass is allowed (0.5—L.O.U.. ).," If the spectral type is larger or smaller by one subclass, a larger range in mass is allowed $0.5-1.0\Msun$ )."
 Here. we assume a stellar mass of 0.75.M...," Here, we assume a stellar mass of $0.75\Msun$."
" An additional model parameter is (he svstem inclination. which we can obtain from the stellar rotation period {πω of 6.76 davs determined from long-term monitoring (Grankin el 22008). the projected stellar rotational velocity c,sin’ (12.1kms|: White Hillenbrand 2004). and the stellar radius /?,. These are related by The stellar radius can be estimated from (he measured stellar Iuminosity aud temperature."," An additional model parameter is the system inclination, which we can obtain from the stellar rotation period $P_{\rm rot}$ of $6.76$ days determined from long-term monitoring (Grankin et 2008), the projected stellar rotational velocity $v_* \sin i$ $12.1 \kms$; White Hillenbrand 2004), and the stellar radius $\Rstar.$ These are related by The stellar radius can be estimated from the measured stellar luminosity and temperature."
 The stellar luminosity. obtained by integrating the stellar component in the SED lit (Fig.," The stellar luminosity, obtained by integrating the stellar component in the SED fit (Fig."
 3: see also Furlan οἱ 22006). is 0.582. at the average Taurus distance of ppc. wilh a range of 0.46—O.STL.. if we account for (he range in distances (126—1173 ppc) derived to individual Taurus objects (Bertout Genova 2006).," 3; see also Furlan et 2006), is $0.58 \Lsun$ at the average Taurus distance of pc, with a range of $0.46 - 0.87\Lsun$ if we account for the range in distances $126-173$ pc) derived to individual Taurus objects (Bertout Genova 2006)."
" For Tag=40002200 IXIN. ,=1.6--. where the uncertainty in the stellar temperature dominates the error on the low luminosity end. and the uncertainty in the distance dominates the error on the high Iuminositv end."," For $T_{\rm eff}=4000 \pm 200$ K, $\Rstar= 1.6^{+0.40}_{-0.22}\Rsun,$ where the uncertainty in the stellar temperature dominates the error on the low luminosity end, and the uncertainty in the distance dominates the error on the high luminosity end."
 An examination of all of the uncertainties in the above properties gives a possible inclination range of 55—90 degrees., An examination of all of the uncertainties in the above properties gives a possible inclination range of $55 - 90$ degrees.
 Since V336 Tau does not have the extreme colors of an edge-on disk svstem (e.g. D'Alessio et 22006). we assume a more modest inclination of 65 degrees for the modeling.," Since V836 Tau does not have the extreme colors of an edge-on disk system (e.g, D'Alessio et 2006), we assume a more modest inclination of 65 degrees for the modeling."
 With these constraints on (he stellar mass and svstem inclination. we can then infer the range of disk radii over which (he emission arises assuming lxeplerian rotation.," With these constraints on the stellar mass and system inclination, we can then infer the range of disk radii over which the emission arises assuming Keplerian rotation."
" The maximum velocity extent of the e—10 CO emission lines (e,90+l0knmsI: see sec,"," The maximum velocity extent of the $v$ =1–0 CO emission lines $v_{\rm max} \sim 90\,\pm 10\kms$; see sec."
" 3.2) and (he inclination range of /=55—90 degrees implies an inner radius of AH,=0.05—0.09 AU for the emitting gas.", 3.2) and the inclination range of $i=55-90$ degrees implies an inner radius of $\Rin = 0.05-0.09$ AU for the emitting gas.
" This range of Hj, is consistent with the values of H4, measured [or Classical T Tauri stars using CO emission line profiles (Najita et 220072: Cary 2007).", This range of $\Rin$ is consistent with the values of $\Rin$ measured for classical T Tauri stars using CO emission line profiles (Najita et 2007a; Carr 2007).
 The overall shape of the emission line profile further suggests that (he emission extends out (o a, The overall shape of the emission line profile further suggests that the emission extends out to a
such large uncertainties emiposecl by several individual motions appear to be unimportant.,such large uncertainties emposed by several individual motions appear to be unimportant.
 But still. for the sake of the statistical completeness. the missing standard. errors of 15 spectroscopic parallaxes had to be completed.," But still, for the sake of the statistical completeness, the missing standard errors of 15 spectroscopic parallaxes had to be completed."
 Sparke Gallagher (2000) state that if the interstellar absorption and the reddening do not introduce problems. the Iuminosities of the main-sequence stars can often be found to within10'4.. leacing to uncertainties in their distance.," Sparke Gallagher (2000) state that if the interstellar absorption and the reddening do not introduce problems, the luminosities of the main-sequence stars can often be found to within, leading to uncertainties in their distance."
 The giant branch is almost. vertical. thus the best. hope for determining a luminosity is within 0.5 in the absolute magnitude. and hence the distance to4.," The giant branch is almost vertical, thus the best hope for determining a luminosity is within 0.5 in the absolute magnitude, and hence the distance to."
. Being in the safe side. the sub giants were assumed to be as giants. thus uncertainty were assigned Lor the missing standard errors of cight giants and four sub giants.," Being in the safe side, the sub giants were assumed to be as giants, thus uncertainty were assigned for the missing standard errors of eight giants and four sub giants."
 The missing standard errors of three svstems (LAL Vir. HIP Aur. HZ Com) with cart components were assigned with a uncertainty as Sparke Gallagher (2000) suggest.," The missing standard errors of three systems (IM Vir, HP Aur, HZ Com) with dwarf components were assigned with a uncertainty as Sparke Gallagher (2000) suggest."
 With a median cistance of 98 xc. the current CAB sample contains the nearby. systems hat the interstellar absorption anc the reddening could »: ignored.," With a median distance of 98 pc, the current CAB sample contains the nearby systems that the interstellar absorption and the reddening could be ignored."
 Moreover. the CAB are popular that they are usually well studied systems that we are confident to apply he rules of Sparke Callagher (2000) for estimating the missing stancard errors of 15 in the list) spectroscopic xwallaxes.," Moreover, the CAB are popular that they are usually well studied systems that we are confident to apply the rules of Sparke Gallagher (2000) for estimating the missing standard errors of 15 in the list) spectroscopic parallaxes."
 LAL Vir and WZ Com are within 60 pe., IM Vir and HZ Com are within 60 pc.
 Thus. with a 287 pe distance. only the error of LIP Aur could of doubted.," Thus, with a 287 pc distance, only the error of HP Aur could be doubted."
 Nevertheless. I0 will not ellect the statistics of the whole sample.," Nevertheless, It will not effect the statistics of the whole sample."
" Phe average propagated errors. at U.V.W for these 15 svstemis were computed as 90= 5.49. OV=+455 and 83V=x31 km/s. which are smaller than the propagated: errors of nine svstems with mz/z20.5. but bigger than the average propagated errors of the whole sample: δὲ=£3.43. 09V=£2.92 and OW=+242 km/s. Finally. after. Gillinge in the missingὃν information in Table 1. the average standard errors on the proper motion components are 0.62 mas/vr in [5n,.cosó and 0.43 mas/vr in ps and the average relative uncertainty of the parallaxes (σε ἐπ) 1s LET."," The average propagated errors at U,V,W for these 15 systems were computed as $\delta U=\pm 5.49$ , $\delta V=\pm 4.55$ and $\delta W=\pm 3.81$ km/s, which are smaller than the propagated errors of nine systems with $\sigma_{\pi}/\pi>0.5$, but bigger than the average propagated errors of the whole sample: $\delta U=\pm 3.43$, $\delta V=\pm 2.92$ and $\delta W=\pm 2.42$ km/s. Finally, after filling in the missing information in Table 1, the average standard errors on the proper motion components are 0.62 mas/yr in $\mu_{\alpha} cos\delta$ and 0.43 mas/yr in $\mu_{\delta}$ and the average relative uncertainty of the parallaxes $\sigma_{\pi}/\pi$ ) is $14.7\%$."
 Unlike the proper motions and the parallaxes. which were mostly taken from the Hipparcos and the Pyveho Catalogs. the radial velocities were collected one by one. from the literature.," Unlike the proper motions and the parallaxes, which were mostly taken from the Hipparcos and the Tycho Catalogs, the radial velocities were collected one by one from the literature."
 Moreover. unlike single stars with a single racial velocity. the binaries ancl the multiple svstems require the racial velocity for the mass center of the svstem (5).," Moreover, unlike single stars with a single radial velocity, the binaries and the multiple systems require the radial velocity for the mass center of the system $\gamma$ )."
 That is. numerous racial velocity measurements are needed. just for computing the orbital parameters together with the velocity of the mass center of a system.," That is, numerous radial velocity measurements are needed just for computing the orbital parameters together with the velocity of the mass center of a system."
 Fortunately. the CAB are popular so that the reliable orbital parameters hack already. been determined for many. svstems.," Fortunately, the CAB are popular so that the reliable orbital parameters had already been determined for many systems."
 However. there are 21 systems in our list (Table 1) which are known to be binaries but. do not. vet possess. determined: orbital elements.," However, there are 21 systems in our list (Table 1) which are known to be binaries but do not yet possess determined orbital elements."
 For such svstems. the mathematical mean of the measured radial velocities was adopted as the center of mass velocity and then the standard. deviation from this mean was taken to be the error estimate.," For such systems, the mathematical mean of the measured radial velocities was adopted as the center of mass velocity and then the standard deviation from this mean was taken to be the error estimate."
 On the other hand. there are many systems with multiple orbit determinations.," On the other hand, there are many systems with multiple orbit determinations."
 Nevertheless. most of the multiple orbit determinations are not independent.," Nevertheless, most of the multiple orbit determinations are not independent."
 That is. the data used. in the previous determination were also used or considered in the later study which gives the most. improved. orbital elements.," That is, the data used in the previous determination were also used or considered in the later study which gives the most improved orbital elements."
 In such cases. it was preferred to use the value of (5) and its associated error. from. the most recently determined orbit unless the most recent study gives unexpectechy large associated errors.," In such cases, it was preferred to use the value of $\gamma$ ) and its associated error from the most recently determined orbit unless the most recent study gives unexpectedly large associated errors."
 Rarely. there are systems with independently determined. orbital parameters.," Rarely, there are systems with independently determined orbital parameters."
 For those. the weighted mean of the systemic velocities (5) and the weighted mean of the associated: errors were used.," For those, the weighted mean of the systemic velocities $\gamma$ ) and the weighted mean of the associated errors were used."
 Those systems are Listed with the multiple. reference. numbers separated by commas after the second semicolon in the last column of Table 1., Those systems are listed with the multiple reference numbers separated by commas after the second semicolon in the last column of Table 1.
 Dillerent authors prefer to. give. dillerent kines of uncertainties associated. with the published: parameters of the orbit., Different authors prefer to give different kinds of uncertainties associated with the published parameters of the orbit.
 In order to maintain consistency. the dillerent types of uncertainties have been transformed into standard errors since most of our data are already expressed. with the standard errors.," In order to maintain consistency, the different types of uncertainties have been transformed into standard errors since most of our data are already expressed with the standard errors."
 Except for the probable. error. the other uncertainties (mean error. standard. error. rms. error and 0) indicate the same confidence level.," Except for the probable error, the other uncertainties (mean error, standard error, rms error and $\sigma$ ) indicate the same confidence level."
 Therefore. they are transferred. cirectly.," Therefore, they are transferred directly."
 However. when transforming the probable errors. (PE) to the standard errors. (SE). the relation of Pleo=0.6755{ was used.," However, when transforming the probable errors (PE) to the standard errors (SE), the relation of $PE = 0.675 SE$ was used."
 Galactic space velocity. components (C. V. M). were computed. together. with their errors. by applying the algorithm. anc the transformation matrices of Johnson Socderblom (1987) to the basic data: celestial coordinates (o. 8). proper motion components (fA. fr). radial velocity (5) and the parallax (x) of each star in Table 1. where the epoch of J2000 coordinates were adopted as described. in the International Celestial Reference System (LCRS) of the llipparcos and the VPvcho Catalogues.," Galactic space velocity components $U$, $V$, $W$ ) were computed together with their errors by applying the algorithm and the transformation matrices of Johnson Soderblom (1987) to the basic data; celestial coordinates $\alpha$, $\delta$ ), proper motion components $\mu_{\alpha}$, $\mu_{\delta}$ ), radial velocity $\gamma$ ) and the parallax $\pi$ ) of each star in Table 1, where the epoch of J2000 coordinates were adopted as described in the International Celestial Reference System (ICRS) of the Hipparcos and the Tycho Catalogues."
 The transformation matrices use the notation of the right handed: svstem., The transformation matrices use the notation of the right handed system.
 Therefore. C. V. M are the components of a velocity. vector of à star with respect. to the Sun. where € is directed toward the Galactic center (=07.5= O°): V is in the direction. of the galactic rotation (/=90°.b O°): and WO ds towards the north Galactic pole (6= 90°).," Therefore, $U$, $V$, $W$ are the components of a velocity vector of a star with respect to the Sun, where $U$ is directed toward the Galactic center $l=0^{o}, b=0^{o}$ ); V is in the direction of the galactic rotation $l=90^{o}, b=0^{o}$ ); and $W$ is towards the north Galactic pole $b=90^{o}$ )."
 The computed. uncertainties are quite small and the averages ave OU)=X343. 01=£2.92 and 98M=42.42 km/s. By inspecting the space velocity vectors (s=vU?|στV2 M7). only Ls 4)) svstems with the uncertainty of the space velocity bigeer than £15 km/s were found.," The computed uncertainties are quite small and the averages are $\delta U=\pm 3.43$, $\delta V=\pm 2.92$ and $\delta W=\pm 2.42$ km/s. By inspecting the space velocity vectors $s=\sqrt{U^{2}+V^{2}+W^{2}}$ ), only 18 ) systems with the uncertainty of the space velocity bigger than $\pm15$ km/s were found."
 Lf those svstenis were removed from the sample. the average uncertainties of the components would reduce to 90=cz24. 90V=x20. and OW=21S km/s. Thus. most of our sample stars have uncertainties very much smaller than the clispersions calculated.," If those systems were removed from the sample, the average uncertainties of the components would reduce to $\delta U=\pm 2.4$, $\delta V=\pm2.0$, and $\delta W=\pm 1.8$ km/s. Thus, most of our sample stars have uncertainties very much smaller than the dispersions calculated."
 Aefore discussing the velocity clispersions and. kinematical implications. it was cecided to inspect the space distribution of the sample CAB.," Before discussing the velocity dispersions and kinematical implications, it was decided to inspect the space distribution of the sample CAB."
 Therefore. the Sun. centeredrectangular galactic coordinates (VOY.Z) corresponding to space velocity components (0.1.M) were caleulated.," Therefore, the Sun centeredrectangular galactic coordinates $(X, Y, Z)$ corresponding to space velocity components $(U, V, W)$ were calculated."
 The computed. coordinates are given in Table 2., The computed coordinates are given in Table 2.
 The projected positions on the galactic plane CX.Y. plane) and on the plane perpenclicular to it CX.Z plane) are displaved in Figure 1.," The projected positions on the galactic plane $X, Y$ plane) and on the plane perpendicular to it $X, Z$ plane) are displayed in Figure 1."
redshifts.,redshifts.
" Therefore, panel B also includes the observed-frame R distribution for optically-bright triples."," Therefore, panel B also includes the observed-frame $R$ distribution for optically-bright triples."
" The distributions indicate that optically-faint triples have R values just as high (actually, even higher) than the optically-bright triples."," The distributions indicate that optically-faint triples have $R$ values just as high (actually, even higher) than the optically-bright triples."
" We conclude that optically-faint triples would not populate the lower-right corner of panel A, and that the absence of such a population represents the true properties of quasars."," We conclude that optically-faint triples would not populate the lower-right corner of panel A, and that the absence of such a population represents the true properties of quasars."
" The absence of quasars in the upper-left corner of panel A (with high R and low Rj)) is also due to intrinsic quasar properties, as demonstrated in panel C. Quasars can be classified as triples only if they have two lobes with Siobe> 1mJy; otherwise, one or both lobes would be undetected."," The absence of quasars in the upper-left corner of panel A (with high $R$ and low ) is also due to intrinsic quasar properties, as demonstrated in panel C. Quasars can be classified as triples only if they have two lobes with $S_\mathrm{lobe}>1$ mJy; otherwise, one or both lobes would be undetected."
" Because of the optical and radio detection limits of our sample, quasars with high R and lowRy,, if they exist, are likely to have (an) undetected lobe(s), and therefore be classified as or rather thantriple."," Because of the optical and radio detection limits of our sample, quasars with high $R$ and low, if they exist, are likely to have (an) undetected lobe(s), and therefore be classified as or rather than."
 Panel C compares the ddistribution of the triples to that of the lobe plus core classes., Panel C compares the distribution of the triples to that of the lobe plus core classes.
" The latter do not show a lower range ofHj,, which indicates that no such population of quasars with high R and low eexists."," The latter do not show a lower range of, which indicates that no such population of quasars with high $R$ and low exists."
" In summary, Figure 4 shows a strong correlation between R andRy."," In summary, Figure \ref{fig:rri_rcl_triple} shows a strong correlation between $R$ and."
". We demonstrated, using distributions in Panels B and C, that the observed correlation is not a selection effect, but is instead an intrinsic property: quasars with high R tend to have highRy,, and vice-versa."," We demonstrated, using distributions in Panels B and C, that the observed correlation is not a selection effect, but is instead an intrinsic property: quasars with high $R$ tend to have high, and vice-versa."
 This observation supports the hypothesis that both parameters are measures of quasar orientation., This observation supports the hypothesis that both parameters are measures of quasar orientation.
 The scatter in the correlation supports the idea that other factors are also influencing these two measurements., The scatter in the correlation supports the idea that other factors are also influencing these two measurements.
 The spectral composite is a valuable tool for studying properties of large samples of astronomical sources., The spectral composite is a valuable tool for studying properties of large samples of astronomical sources.
" In recent years, the advent of large spectroscopic surveys has allowed for increasingly detailed studies of AGN spectral composites (????7).."," In recent years, the advent of large spectroscopic surveys has allowed for increasingly detailed studies of AGN spectral composites \citep{francis91,bakerHunstead95,zheng97,brotherton01,vandenberk01}."
" Spectra of many objects can be combined to create composites with high signal-to-noise, allowing average properties to be compared across sample types."," Spectra of many objects can be combined to create composites with high signal-to-noise, allowing average properties to be compared across sample types."
" Here we compare spectral properties of quasar radio morphology classes, and investigate whether a selection based on radio properties implies concomitant changes in optical spectra."," Here we compare spectral properties of quasar radio morphology classes, and investigate whether a selection based on radio properties implies concomitant changes in optical spectra."
" Creating a spectral composite involves several steps: shifting each input spectrum to its rest frame, rebinning the spectra to matching wavelength grids and resolution, normalizing their fluxes, and stacking into a final composite."," Creating a spectral composite involves several steps: shifting each input spectrum to its rest frame, rebinning the spectra to matching wavelength grids and resolution, normalizing their fluxes, and stacking into a final composite."
 Variations of these steps can lead to significant differences in the resulting composite , Variations of these steps can lead to significant differences in the resulting composite \citep{francis91}.
We redshift-corrected each spectrum using the (?)..S07 redshift (Section 2.2))., We redshift-corrected each spectrum using the S07 redshift (Section \ref{subsec:sdss}) ).
" Spectra were then re-binned, conserving flux, onto a common wavelength grid with bins 69 km s! wide, the same sampling as observed SDSS spectra."," Spectra were then re-binned, conserving flux, onto a common wavelength grid with bins $\sim69$ km $^{-1}$ wide, the same sampling as observed SDSS spectra."
" We used the ? quasar spectral composite as a template for rescaling our input spectra to a common flux level, using the following method."," We used the \citet{vandenberk01} quasar spectral composite as a template for rescaling our input spectra to a common flux level, using the following method."
" In the region of wavelength overlap, we found the flux ratio of the input spectrum to the Vanden Berk et al."," In the region of wavelength overlap, we found the flux ratio of the input spectrum to the Vanden Berk et al."
 composite at each wavelength; however we ignored the region blueward of the Lya emission line (A<1250A)) so as to avoid uncertainties due to redshift-dependent absorption in the Lya forest region., composite at each wavelength; however we ignored the region blueward of the $\alpha$ emission line $\lambda<1250$ ) so as to avoid uncertainties due to redshift-dependent absorption in the $\alpha$ forest region.
" We then determined the median flux ratio in the overlap region, and used this value to rescale the input spectrum."," We then determined the median flux ratio in the overlap region, and used this value to rescale the input spectrum."
 Our resulting composites are in arbitrary flux density units (but on the same scale as the Vanden Berk et al., Our resulting composites are in arbitrary flux density units (but on the same scale as the Vanden Berk et al.
 composite)., composite).
 We stacked the rescaled spectra into two different final composites: a “median” composite and a “geometric mean” composite. ," We stacked the rescaled spectra into two different final composites: a “median"" composite and a “geometric mean"" composite. ("
"the latter, negative flux values were ignored.)","For the latter, negative flux values were ignored.)"
" A (Formedian composite preserves the relative fluxes of emission features, while a geometric mean preserves the overall continuum shape for power-law spectra (?)."," A median composite preserves the relative fluxes of emission features, while a geometric mean preserves the overall continuum shape for power-law spectra \citep{vandenberk01}."
" We note that our composites do not include BAL quasars, as BALs can influence a spectrum's shape."," We note that our composites do not include BAL quasars, as BALs can influence a spectrum's shape."
" We compared the median composite for the 72,223 quasars (both radio and radio-quiet) without BALs to the median composite from Vanden Berk et al.,"," We compared the median composite for the 72,223 quasars (both radio and radio-quiet) without BALs to the median composite from Vanden Berk et al.,"
 which also excludes BAL quasars., which also excludes BAL quasars.
" The two spectra agree very well and are nearly indistinguishable to the eye, except at the extreme wavelength values, where fewer spectra contribute and thus the composites are noisier."," The two spectra agree very well and are nearly indistinguishable to the eye, except at the extreme wavelength values, where fewer spectra contribute and thus the composites are noisier."
 The spectral composites for the entire sample and for individual morphological classes are discussed in the remainder of this section., The spectral composites for the entire sample and for individual morphological classes are discussed in the remainder of this section.
" Table 5 presents the composite for the radio-quiet quasars, for the full sample of radio quasars, and for the radio morphology classes."," Table \ref{table:composites_main} presents the composite for the radio-quiet quasars, for the full sample of radio quasars, and for the radio morphology classes."
 Individual quasar spectra have different rest-frame wavelength coverage owing to their wide redshift range, Individual quasar spectra have different rest-frame wavelength coverage owing to their wide redshift range
We consider the photometric studies of PAIS stars as very naportaut for their exact classification.,We consider the photometric studies of PMS stars as very important for their exact classification.
 The xoblenis fro superposition of diftorent tvpes of variability cai be solved bv. collectiie of long-term photometric data from the photographic plate archives and frou phoonmoetrie monitoring in t1ο present time., The problems from superposition of different types of variability can be solved by collecting of long-term photometric data from the photographic plate archives and from photometric monitoring in the present time.
 Another dispiited poiut that can be solved by both photometric aud spectral monitoring is searching for correlation beween the brightucss of he star aud the equivalent width of Ta line., Another disputed point that can be solved by both photometric and spectral monitoring is searching for correlation between the brightness of the star and the equivalent width of $\alpha$ line.
 Such a correlation between he deep iiwus in brightucss aud he increasing of he equivalent width of Wa line would be confirmation or the UNor ype of variability of GAL Cep., Such a correlation between the deep minimums in brightness and the increasing of the equivalent width of $\alpha$ line would be confirmation for the UXor type of variability of GM Cep.
 This work was partly supported by erauts DO -s5 aud DC) 02-273 of the Nationsd Science Fined of the Miuistrv of Education. Youth aud Scieuce. Bulgaria.," This work was partly supported by grants DO 02-85 and DO 02-273 of the National Science Fund of the Ministry of Education, Youth and Science, Bulgaria."
 The authors thank the Director of Skiualsas Observatory Prof. L Papamastorakis and Prof. 1. Papadakis for the telescope time., The authors thank the Director of Skinakas Observatory Prof. I. Papamastorakis and Prof. I. Papadakis for the telescope time.
 This research has uade use of the NASA Astrophysics Data Syste., This research has made use of the NASA Astrophysics Data System.
near-infrared.,near-infrared.
 Fanetal.(2006) determined new ages for 91 MOI GCs from Jiangetal.(2003) based on improved photometric data aud BCO3 models.," \citet{fan06}
 determined new ages for 91 M31 GCs from \citet{jiang03} based on improved photometric data and BC03 models."
 Iu particular. Maetal.(2007) derived the age of Dy comparing its photometric data with BCO3 models.," In particular, \citet{Ma07} derived the age of by comparing its photometric data with BC03 models."
 The age obtained by Maetal.(2007) is 9.5!mi Cyr. which is consisteut with the determination of 10!2 Cyr by Brown(2001a) using the maim-sequence photometry.," The age obtained by \citet{Ma07} is $9.5_{-0.99}^{+1.15}$ Gyr, which is consistent with the determination of $10_{-1}^{+2.5}$ Gyr by \citet{brown04a} using the main-sequence photometry."
 The nearest large GC system outside the Milly Way is that of the Andromeda Galaxy (= M31). which is located at a distance of 770 kpe (Freecanan&Aacore1990).," The nearest large GC system outside the Milky Way is that of the Andromeda Galaxy (= M31), which is located at a distance of 770 kpc \citep{Freedman90}."
. The first M31. GC resolved iuto stars was studied from the eround by Teasleyetal.(L988).. who only resolved the red giant brauch of G1.," The first M31 GC resolved into stars was studied from the ground by \citet{Heasley88}, who only resolved the red giant branch of G1."
 Subsequently. some authors such as Ajharetal.(1996).. FusiPeccictal. (1996). Richetal. (1996).. Hollandetal. (1997)... al. (2000)... Williams&IIodge (2001a).. aud Williams&Ποάσο(2001b) have used images frou the HST/WEDPC? to construct the CAIDs of M21 star clusters m order to determine their metalliities. reddening values. aud ages.," Subsequently, some authors such as \citet{Ajhar96}, \citet{fusi96}, \citet{rich96}, \citet{hfr97}, \citet{Jablonka00}, \citet{Williams01a}, and \citet{Williams01b} have used images from the /WFPC2 to construct the CMDs of M31 star clusters in order to determine their metallicities, reddening values, and ages."
 However. these CMDs are not deep euough to show conspicuous main-sequence turi-offs.," However, these CMDs are not deep enough to show conspicuous main-sequence turn-offs."
 Since the Iuninositv of the horizoutal brauch (IB) 1- stellar populations older than about 8 Cyr is expected to be independent of age aud ouly mildly dependent o- metallicity. it is widely used a distance indicator (seeGallartetal.2005.audreferences therein).," Since the luminosity of the horizontal branch (HB) in stellar populations older than about 8 Gyr is expected to be independent of age and only mildly dependent on metallicity, it is widely used a distance indicator \citep[see][and references therein]{Gallart05}."
 In addition. IE stars are fundamental standard candles for Population IL systems. and consequently are important tools for deteriiuiug ages of GCs from the main-sequence turn-off hinunosity (Richetal.1996).," In addition, HB stars are fundamental standard candles for Population II systems, and consequently are important tools for determining ages of GCs from the main-sequence turn-off luminosity \citep{rich96}."
. Two of the first studies about the WB for M31. GCs were that of Richetal.(1996) and. FusiPeccietal. (1996)., Two of the first studies about the HB for M31 GCs were that of \citet{rich96} and \citet{fusi96}.
.. They used the observed data from the Wide Field Planetary Camera-2 (WFPC2) aud Faint Object Camera (FOC) onboard theHST to make the first tentative detection of WBin Gl. Boos. BOLS. D225. D313. 3358. D105. aud D168. deriving the appareut magnitude of the IIB for these GCs to be in the range 25.29<Vox 25.66.," They used the observed data from the Wide Field Planetary Camera-2 (WFPC2) and Faint Object Camera (FOC) onboard the to make the first tentative detection of HB in G1, B006, B045, B225, B343, B358, B405, and B468, deriving the apparent magnitude of the HB for these GCs to be in the range $25.29<V<25.66$ ."
 In addition. FusiPeccietal.(1996). firstly preseuted a direct calibration. for the mean absolute magnitude of the IIB at the instability strip with varviug metallicity for NI. CC.," In addition, \citet{fusi96} firstly presented a direct calibration for the mean absolute magnitude of the HB at the instability strip with varying metallicity for M31 GCs."
 B379 was firstly detected by Sharoy(1973). (N0.19). and confirmed by Sargeutetal.(1977) (No.312=SKB 312. which was used by Brownetal.(20014): or =S312. which was used by Maetal. (2007)..)," B379 was firstly detected by \citet{sh73} (No.19), and confirmed by \citet{sarg77} $\rm
No.312=SKHB~312$ , which was used by \citet{brown04a}; or $=\rm S312$, which was used by \citet{Ma07}. .)"
 aud Dattistinietal.(1987). (No.379=B379. which is used in the RBC (Calleti ct al.," and \citet{battis87} $\rm No.379=B379$, which is used in the RBC (Galleti et al."
 2001. 2006. 2007) aud his paper.).," 2004, 2006, 2007) and this paper.)."
 B379 is located in the halo of M31. at a sojected distance of about 59° (=13 kpe) from the ealaxvs nucleus.," B379 is located in the halo of M31, at a projected distance of about $59\arcmin$ $\rm =13~kpc$ ) from the galaxy's nucleus."
 It is a fact that B379 is a common malo GC. however. it ix among the first extragalactie GCs whose age was accurately estimateck dy iiadnu-sequeuce photometry (Brownetal.20012). based on its CMD from he extremely deep images observed with the HST/ACS.," It is a fact that B379 is a common halo GC, however, it is among the first extragalactic GCs whose age was accurately estimated by main-sequence photometry \citep{brown04a} based on its CMD from the extremely deep images observed with the /ACS."
 Iu this paper. we re-cdetermine the age. metallicity. reddening value and distance modulus for D379 bv comparing its CMD constructed by Brownetal.(2001a) with isochrones of the Padova stellar evolutionary models.," In this paper, we re-determine the age, metallicity, reddening value and distance modulus for B379 by comparing its CMD constructed by \citet{brown04a} with isochrones of the Padova stellar evolutionary models."
 The paper is organized as follows., The paper is organized as follows.
 In 83. we deseribe the results of photometric data based ou the IIST/CS observations for D379.," In 3, we describe the results of photometric data based on the /ACS observations for B379."
 In 81. we coustrain the age. metallicity. reddening value. aud distance modulus for D379.," In 4, we constrain the age, metallicity, reddening value, and distance modulus for B379."
 At last. we will eive a σον in 5.," At last, we will give a summery in 5."
 Brownetal(2001a) used the imidn-sequeuce photometry to determine the age of D379 based on the CAD constructed using the extremely deep images from he HST/ACS observations., \citet{brown04a} used the main-sequence photometry to determine the age of B379 based on the CMD constructed using the extremely deep images from the /ACS observations.
 This CMD reached more han 1.5 mag below the maiu-sequence turn-off. aud firstly allowed a direct age estimate from the turn-off for an extragalactic cluster.," This CMD reached more than 1.5 mag below the main-sequence turn-off, and firstly allowed a direct age estimate from the turn-off for an extragalactic cluster."
 By comparison to isochrones of VandeuBereetal.(2006).. Brownetal.(20012). derived he age of D379 to be 10!2 Car.," By comparison to isochrones of \citet{vandenberg06}, \citet{brown04a} derived the age of B379 to be $10_{-1}^{+2.5}$ Gyr."
 Maetal.(2007) determined the age of D379 by comparing its inulti-color photometric data which including the near-ultraviolet (NUV) from the Nearby Galaxies. Survey (NCS) of heExplorer (GALEN) (Revctal.2005.2007).. broad-band CBVR (Battistinietal.1987:Reed.Harris&199 0).. 9 BATC intermediate-band filters aud Two Micron All Sky Survey (2MASS) ΗΝ... with SSP models of BCO3.," \citet{Ma07} determined the age of B379 by comparing its multi-color photometric data which including the near-ultraviolet (NUV) from the Nearby Galaxies Survey (NGS) of the ) \citep{rey05,rey06}, broad-band $UBVR$ \citep{battis87,rhh94}, 9 BATC intermediate-band filters and Two Micron All Sky Survey (2MASS) $JHK_s$, with SSP models of BC03."
 These photometric data constitute the SEDs of B379 covering 2267 20000A., These photometric data constitute the SEDs of B379 covering $2267-20000$.
. The age of B379 determined by Maetal.(2007) 1 9.51083 Gyr. which is cousisteut with the determination of 10!2 Gyr by Brownetal.(200la).," The age of B379 determined by \citet{Ma07}
 is $9.5_{-0.99}^{+1.15}$ Gyr, which is consistent with the determination of $10_{-1}^{+2.5}$ Gyr by \citet{brown04a}."
. However. Ός00 models are based on the Padova evolutionary tracks.," However, BC03 models are based on the Padova evolutionary tracks."
 So. if is necessary to compare the age scale between the Padova evolutionary tracks and the Victoria-Reeina isochrones used im Brownetal.(200la).. aud only if these two evolutionary tracks have the consisteut age scale for D379. the results” comparison between Brownetal.(200la) aud. Maotal.(2007) is imcanineful.," So, it is necessary to compare the age scale between the Padova evolutionary tracks and the Victoria-Regina isochrones used in \citet{brown04a}, and only if these two evolutionary tracks have the consistent age scale for B379, the results' comparison between \citet{brown04a} and \citet{Ma07} is meaningful."
 As an cxample. Maetal.(2007) drew the isoclirones with 10 Car and the solar metallicity. and found that the matching is very eood iu the maim-sequence (MS) and the subeiaut branch (SCB} (seeMaetal.2007.for details}...," As an example, \citet{Ma07} drew the isochrones with 10 Gyr and the solar metallicity, and found that the matching is very good in the main-sequence (MS) and the subgiant branch (SGB) \citep[see][for details]{Ma07}."
 Towever. we should check whether the age of D379 cau be estimated to be ~10 Cxr based on the Padova evolutionary tracks.," However, we should check whether the age of B379 can be estimated to be $\sim 10$ Gyr based on the Padova evolutionary tracks."
 This is one of the key contributions of the present paper., This is one of the key contributions of the present paper.
 The observed data of D379 in this study are from Brownetal.(200[3).. who constructed the CMD of B379 using the images observed with the ACS observations iu the F606W and the FalAW filters.," The observed data of B379 in this study are from \citet{brown04a}, who constructed the CMD of B379 using the images observed with the ACS observations in the F606W and the F814W filters."
 Using the ACS Wide-Field Camera (WFC). Brownetal.(2003) obtained deep optical nuages of a field. 51 from the uncleus ou the southeast münor axis of the M31 halo including D379. which are 39.1 hr iu the FOOGW filter aud. 15.L lr in the FSIIW filter.," Using the ACS Wide-Field Camera (WFC), \citet{brown03} obtained deep optical images of a field, $51'$ from the nucleus on the southeast minor axis of the M31 halo including B379, which are 39.1 hr in the F606W filter and 45.4 hr in the F814W filter."
 Brownetal.(2001a). preseuted the CMD of B379 based on these ACS observations., \citet{brown04a} presented the CMD of B379 based on these ACS observations.
 The resulting CMD reached i0372230.5 nag. which is the first CMD of extragalactic clusters reaching more than 1.5 mag below the main-sequence turn-off.," The resulting CMD reached $m_{V}\approx 30.5$ mag, which is the first CMD of extragalactic clusters reaching more than 1.5 mag below the main-sequence turn-off."
 These observations firstly allow a direct age estimate from the turn-off for an extragalactic cluster., These observations firstly allow a direct age estimate from the turn-off for an extragalactic cluster.
 By comparison to isochrones of VandeuBereetal. (2006).. Brownetal.(20012). derived the age of D379 to be 10!2 Cyr.," By comparison to isochrones of \citet{vandenberg06}, , \citet{brown04a} derived the age of B379 to be $10_{-1}^{+2.5}$ Gyr."
 In Brownctal. la)... the CMD of D379 was constructed. from stars," In \citet{brown04a}, , the CMD of B379 was constructed from stars"
an additional constraint everv model of SN 90016] needs to salisly.,an additional constraint every model of SN 2001el needs to satisfy.
 We will sketch here and critique several possibilities., We will sketch here and critique several possibilities.
 In order to make a high-velocity shell of calcium-rich material modestly optically. thick so that il can obscure a patel of the photosphere aud induce a significant. polarization feature. (here must be a substantial mass of high-velocity caleium.," In order to make a high-velocity shell of calcium-rich material modestly optically thick so that it can obscure a patch of the photosphere and induce a significant polarization feature, there must be a substantial mass of high-velocity calcium."
 A simple static atmosphere argunient suggests the following., A simple static atmosphere argument suggests the following.
 Taking an effective opacity of order | em? ! and shell radius and thickness of order 104? cm. the mass of calcium must be of order 5xLO7frHr. (lo give an optical depth of order unity. where f is the filling factor of the high-velocity shell.," Taking an effective opacity of order 1 $^2$ $^{-1}$ and shell radius and thickness of order $10^{15}$ cm, the mass of calcium must be of order $5\times10^{-3}~f\kappa^{-1}R_{15}^2$ to give an optical depth of order unity, where $f$ is the filling factor of the high-velocity shell."
 A more rigorous approach is to consider the Sobolev optical depth of the line: where οναν = (in a homologously expanding atmosphere and y is the integrated line opacily given by where { and νι are the oscillator strength ancl frequency [or Cie line (hat represents a transition form (he lower (1) to upper (1) levels., A more rigorous approach is to consider the Sobolev optical depth of the line: where dr/dv = t in a homologously expanding atmosphere and $\chi$ is the integrated line opacity given by where $f_{ul}$ and $\nu_{ul}$ are the oscillator strength and frequency for the line that represents a transition form the lower (l) to upper (u) levels.
" Neelectüng the second term in parentlesis involving the statistical weights and writing the population in the lower level as €xp(—Ft). where neg is the total munber of calcium atoms. the Sobolev optica depth can be written as: where E, is the energv of the lower level."," Neglecting the second term in parenthesis involving the statistical weights and writing the population in the lower level as $n_l = n_{Ca}~exp(-E_l/kt)$ , where $n_{Ca}$ is the total number of calcium atoms, the Sobolev optical depth can be written as: where $E_l$ is the energy of the lower level."
" With f,;3. E,=60.533 ! — L2xLo! ere. and a temperature of about 5000 Ix. (he Sobolev optical depth can be written as: where /4 is the time in units of 10"" s. The mass of caleum in a spherical shell with filling factor. F. ean thus be written:"," With $f_{ul} \sim 3$, $E_l = 60,533$ $^{-1}$ = $1.2\times10^{-11}$ erg, and a temperature of about 5000 K, the Sobolev optical depth can be written as: where $t_6$ is the time in units of $10^6$ s. The mass of calcium in a spherical shell with filling factor, f, can thus be written:"
solutions for the Hines and the emission wings using RV measurements on the DAO out-of-eclipse spectra.,solutions for the lines and the emission wings using RV measurements on the DAO out-of-eclipse spectra.
 We used the FOTEL program (?) with input parameters from the solution of? and only allowed a convergency of the periastron epoch T and the velocity semi-amplitude Kj., We used the FOTEL program \citep{Hadrava2004} with input parameters from the solution of \citet{Chadima2010} and only allowed a convergency of the periastron epoch $T$ and the velocity semi-amplitude $K_1$.
 Results are summarized in Table [.., Results are summarized in Table \ref{tab3}.
 In Fig. 5..," In Fig. \ref{orbit},"
 the orbital solution found trom the lines alone and the solution from the wings of are compared with the solution of ?.., the orbital solution found from the lines alone and the solution from the wings of are compared with the solution of \citet{Chadima2010}.
" It is seen that the emission roughly follows the orbital motion of the Imes, Le.. the primary. and therefore must be associated with this component."," It is seen that the emission roughly follows the orbital motion of the lines, i.e., the primary, and therefore must be associated with this component."
 We infer that a small difference between both solutions ts mainly caused by an insufficient interval covered by data used for both orbital solutions (when compared with the orbital period of the Aur))., We infer that a small difference between both solutions is mainly caused by an insufficient interval covered by data used for both orbital solutions (when compared with the orbital period of the ).
 To verity that the emission. originates near. the primary. we selected two rather symmetric profiles trom substantially different orbital phases and overplotted them on a heliocentric wavelength scale in Fig. 6..," To verify that the emission originates near the primary, we selected two rather symmetric profiles from substantially different orbital phases and overplotted them on a heliocentric wavelength scale in Fig. \ref{halfa}."
 One can see that the whole profile. including the double emission wings. undergoes the RV shift corresponding to the motion ofstar in orbit.," One can see that the whole profile, including the double emission wings, undergoes the RV shift corresponding to the motion of in orbit."
 This behavior provides strong evidence that the emission is associated with the F-type primary., This behavior provides strong evidence that the emission is associated with the F-type primary.
 During eclipse. undergoes prominent line. profile changes.," During eclipse, undergoes prominent line profile changes."
 Figure 7 shows a remarkable fact: the CI of the absorption core steadily decreases with the approaching eclipse., Figure \ref{Ha_orig} shows a remarkable fact: the CI of the absorption core steadily decreases with the approaching eclipse.
 An important finding is that this decreaseeclipse. thus confirming the behavior first reported by ?..," An important finding is that this decrease, thus confirming the behavior first reported by \citet{Kuiper1937}."
 The EW increased during the same time period., The EW increased during the same time period.
 The RV of the central absorption core increased significantly during the first half of the eclipse. but after mideclipse. then suddenly reverted back to negative values.," The RV of the central absorption core increased significantly during the first half of the eclipse, but after mideclipse, then suddenly reverted back to negative values."
 The emission peaks disappear during eclipse., The emission peaks disappear during eclipse.
" For the current eclipse, the red emission peak had disappeared by HID 2455259 and the blue peak had vanished by HJD 2455329."," For the current eclipse, the red emission peak had disappeared by HJD 2455259 and the blue peak had vanished by HJD 2455329."
 The red emission peak subsequently reappeared after HJD 2455524., The red emission peak subsequently reappeared after HJD 2455524.
 From a study of the line profile during the last eclipse (?).. We expect the evolution of the profile in the second half of the current eclipse to follow the sequence of events described above in reverse.," From a study of the line profile during the last eclipse \citep{Barsony1986}, we expect the evolution of the profile in the second half of the current eclipse to follow the sequence of events described above in reverse."
" The observations should show a gradual increase in the absorption core CI. accompanied by a decrease in its EW, and a gradual return of the absorption core RV to the out-of-eclipse values."," The observations should show a gradual increase in the absorption core CI, accompanied by a decrease in its EW, and a gradual return of the absorption core RV to the out-of-eclipse values."
 It is expected that both emission peaks will reappear in reverse order to what was seen in the first half of the eclipse., It is expected that both emission peaks will reappear in reverse order to what was seen in the first half of the eclipse.
 Based on the results of previous studies of the shell spectrum (discussed in Sect. ??)).," Based on the results of previous studies of the shell spectrum (discussed in Sect. \ref{intro}) ),"
 we adopted the ? interpretation of the behavior.," we adopted the \citet{Kuiper1937}
 interpretation of the behavior."
" We assume that (apart trom any intrinsic line profile changes, as seen for other lines) changes in the profile during eclipse are due to additional absorption of the primary’s light by the outer. optically thin atmosphere of the companion's occulting disk."," We assume that (apart from any intrinsic line profile changes, as seen for other lines) changes in the profile during eclipse are due to additional absorption of the primary's light by the outer, optically thin atmosphere of the companion's occulting disk."
"predicts the frequency of the most unstable modes for each /, but significantly overestimates that of other modes.","predicts the frequency of the most unstable modes for each $l$, but significantly overestimates that of other modes."
 This is consistent with the result of Section 3 that the acoustic radial time is significantly shorter than that extracted from the eigenspectrum., This is consistent with the result of Section \ref{sec:radial_time} that the acoustic radial time is significantly shorter than that extracted from the eigenspectrum.
" It is also consistent with the fact that the SASI eigenspectrum is much more densely populated than the acoustic eigenspectrum (compare Figure 9 and Figure 10,, remembering that the frequency scale is different)."," It is also consistent with the fact that the SASI eigenspectrum is much more densely populated than the acoustic eigenspectrum (compare Figure \ref{fig:acoustic_spectrum} and Figure \ref{fig:SASI_spectrum}, remembering that the frequency scale is different)."
" Similarly to the toy model, the most unstable SASI modes for each spherical harmonics index | are close to the purely acoustic branch n,=0."," Similarly to the toy model, the most unstable SASI modes for each spherical harmonics index $l$ are close to the purely acoustic branch $n_r=0$."
" This can again be interpreted by a higher efficiency of the advective-acoustic cycle, as well as a constructive interference with the purely acoustic mode."," This can again be interpreted by a higher efficiency of the advective-acoustic cycle, as well as a constructive interference with the purely acoustic mode."
 It is also probably linked with the observation of ? that the growth rate of a mode is maximum when the advection time coincides with the azimuthal acoustic time., It is also probably linked with the observation of \citet{fernandez09a} that the growth rate of a mode is maximum when the advection time coincides with the azimuthal acoustic time.
" Finally, we observe that the (stable) radial modes are systematically shifted with respect the theoretical curve."," Finally, we observe that the (stable) radial modes are systematically shifted with respect the theoretical curve."
" This can be explained by a phase shifttaking place during the advective-acoustic cycle, as shown by ?.."," This can be explained by a phase shifttaking place during the advective-acoustic cycle, as shown by \citet{fernandez09b}."
" Indeed, a phase shift of ᾧ=π then gives the following frequency spectrum for the |=0 modes: which is in much better agreement with the numerical eigenmodes."," Indeed, a phase shift of $\Phi=\pi$ then gives the following frequency spectrum for the $l=0$ modes: which is in much better agreement with the numerical eigenmodes."
" 'Two new arguments have been presented which allow us to conclude that SASI cannot be explained by a purely acoustic mechanism, and should thus be interpreted as an advective-acoustic cycle."," Two new arguments have been presented which allow us to conclude that SASI cannot be explained by a purely acoustic mechanism, and should thus be interpreted as an advective-acoustic cycle."
" Contrary to previous works, these two arguments do not rely on a WKB analysis of the acoustic waves, and are therefore valid for realistic values of the shock radius."," Contrary to previous works, these two arguments do not rely on a WKB analysis of the acoustic waves, and are therefore valid for realistic values of the shock radius."
 First we proposed a method to extract a radial propagation timescale from the frequency difference between two successive SASI modes., First we proposed a method to extract a radial propagation timescale from the frequency difference between two successive SASI modes.
" This timescale agrees fairly well with the advective-acoustic time, but strongly disagrees with the radial acoustic time."," This timescale agrees fairly well with the advective-acoustic time, but strongly disagrees with the radial acoustic time."
" This discards the purely acoustic cycle as the driving mechanism of SASI, and argues in favour of the advective-acoustic cycle."," This discards the purely acoustic cycle as the driving mechanism of SASI, and argues in favour of the advective-acoustic cycle."
 A detailed analysis of the frequency spectrum provided in Section 5 confirmed the validity of this argument., A detailed analysis of the frequency spectrum provided in Section \ref{sec:details} confirmed the validity of this argument.
" Second, we described a new method to compute purely acoustic modes by artificially subtracting the advected wave created by the shock."," Second, we described a new method to compute purely acoustic modes by artificially subtracting the advected wave created by the shock."
" All these acoustic modes were found to be stable, thus showing that the advected entropy-vorticity wave created by the oscillation of the shock plays a central role in the instability mechanism."," All these acoustic modes were found to be stable, thus showing that the advected entropy-vorticity wave created by the oscillation of the shock plays a central role in the instability mechanism."
" As a side product, we provided an analytical description of the frequency spectra of the advective- and purely acoustic cycles."," As a side product, we provided an analytical description of the frequency spectra of the advective-acoustic and purely acoustic cycles."
 This description was checked in the planar toy model of ? and the model of ?.., This description was checked in the planar toy model of \citet{foglizzo09} and the model of \citet{blondin06}.
" The similarity between the frequency spectra in these two models, supports the relevance of the toy model for core collapse supernovae."," The similarity between the frequency spectra in these two models, supports the relevance of the toy model for core collapse supernovae."
" Finally, we note that unexplained phase shifts during the coupling processes were observed in the second model."," Finally, we note that unexplained phase shifts during the coupling processes were observed in the second model."
 This calls for a better understanding of the coupling process in a cooling layer., This calls for a better understanding of the coupling process in a cooling layer.
 The authors are grateful to John Blondin for stimulating discussions., The authors are grateful to John Blondin for stimulating discussions.
 This work has been partially funded by the Vortexplosion project ANR-06-JCJC-0119 and the SN2NS project ANR-10-BLAN-0503., This work has been partially funded by the Vortexplosion project ANR-06-JCJC-0119 and the SN2NS project ANR-10-BLAN-0503.
 J.G. acknowledges support by STFC., J.G. acknowledges support by STFC.
" Following ?,, we adopt the variables 05,0q,0f,óh defined by:"," Following \citet{yamasaki08}, , we adopt the variables $\delta S,\delta q,\delta f,\delta h$ defined by:"
"where Ε]γαοὓς is the observable Lyo flux per unit of star formation, 1s the optical depth due to dust in the galaxy, fesc is the fraction of ionizing photons that escape the star-forming galaxy and thus create noLya photons at the source, and fray is the total fraction of Lya photons that penetrate the intergalactic medium and reach us.","where ${\rm F_{Ly\alpha,obs}}$ is the observable $\alpha$ flux per unit of star formation, is the optical depth due to dust in the galaxy, ${\rm f_{esc}}$ is the fraction of ionizing photons that escape the star-forming galaxy and thus create no$\alpha$ photons at the source, and ${\rm f_{IGM}}$ is the total fraction of $\alpha$ photons that penetrate the intergalactic medium and reach us."
" Fryqem is the maximum, ionization-bounded flux of Lya photons produced per unit of star formation; it depends on the IMF and metallicity."," ${\rm F_{Ly\alpha,em}}$ is the maximum, ionization-bounded flux of $\alpha$ photons produced per unit of star formation; it depends on the IMF and metallicity."
 Recent models and the of high-redshift galaxies provide some constraints interpretationon these quantities., Recent models and the interpretation of high-redshift galaxies provide some constraints on these quantities.
" Theoretical models predict a top-heavy IMF for the first generation of stars formed in metal-free gas (Bromm,Coppi,&Larson1999,2002;Abel,Norman 2000)."," Theoretical models predict a top-heavy IMF for the first generation of stars formed in metal-free gas \citep{Bro99,Bro02,Abe00}."
". These stars can radiate more than anBryan, order of magnitude more ionizing photons per unit solar mass than a Salpeter IMF, resulting in much stronger Lya and II((A1640) (Bromm,Kudritzki,&Loeb2001;Schaerer 2003)."," These stars can radiate more than an order of magnitude more ionizing photons per unit solar mass than a Salpeter IMF, resulting in much stronger $\alpha$ and $\lambda1640$ ) \citep{Bro01,Sch03}."
". However, only the initial generations of stars will be formed from completely pristine gas; thus, one major source of uncertainty is whether luminous bursts of star formation later in the process (z=8) will be too enriched to exhibit extremely top-heavy IMFs."," However, only the initial generations of stars will be formed from completely pristine gas; thus, one major source of uncertainty is whether luminous bursts of star formation later in the process $z=8$ ) will be too enriched to exhibit extremely top-heavy IMFs."
" The factor e""ws(1—fesc) is the fraction of emitted ionizing photons available for conversion to Lya; it includes the of and Lya photons from dust near the newly absorptionforming stars.", The factor ${\rm e^{-\tau_{dust}}\ (1-f_{esc})}$ is the fraction of emitted ionizing photons available for conversion to $\alpha$; it includes the absorption of ionizing and $\alpha$ photons from dust near the newly forming stars.
" ionizingThere are, at present, few direct observational constraints on this factor at any redshift."," There are, at present, few direct observational constraints on this factor at any redshift."
" At z»3, Pettinietal.(1998) estimate est~0.16—0.4."," At $z \sim 3$, \citet{Pet98} estimate ${\rm e}^{-\tau_{\rm dust}}
\sim 0.16 - 0.4$ ."
" But the stellar populations at even higher redshift are probably younger, less chemically evolved, and hence less dusty."," But the stellar populations at even higher redshift are probably younger, less chemically evolved, and hence less dusty."
" For the escape fraction, Steidel,Pettini,&Adelberger(2001) combine z~3 Lyman-break galaxy spectra to estimate an average f»2,0.07—0.1."," For the escape fraction, \citet{Ste01} combine $z \sim 3$ Lyman-break galaxy spectra to estimate an ${\rm f_{esc} \gtrsim
0.07-0.1}$."
 The discovery of Gunn-Peterson troughs in the spectra of quasars at z=6 indicate that the average neutral fraction of the medium exceeds xy~ at z>6 (Beckeretintergalactical.2001;Fan2002).," The discovery of Gunn-Peterson troughs in the spectra of quasars at $z \gtrsim 6$ indicate that the average neutral fraction of the intergalactic medium exceeds $x_{\rm HI} \sim 10^{-2}$ at $z
\gtrsim 6$ \citep{Bec01,Fan02}."
". At face value, this result suggests that Lya emission from beyond this redshift will not penetrate the IGM."," At face value, this result suggests that $\alpha$ emission from beyond this redshift will not penetrate the IGM."
" However, Haiman(2002) shows that a substantial fraction of the emittedLya photons escape scattering by the IGM when a galaxy ionizes a local bubble in its immediate surroundings."," However, \citet{Hai02} shows that a substantial fraction of the emitted$\alpha$ photons escape scattering by the IGM when a galaxy ionizes a local bubble in its immediate surroundings."
" Extending this idea using dynamical models of the IGM and galactic winds, Santos(2003) derives values of fray ranging from <0.002 to ~1."," Extending this idea using dynamical models of the IGM and galactic winds, \citet{San03} derives values of $_{\rm IGM}$ ranging from $\lesssim 0.002$ to $\sim 1$."
 This broad range illustrates the substantial uncertainties involved in estimating the expected Lyo flux., This broad range illustrates the substantial uncertainties involved in estimating the expected $\alpha$ flux.
" In general, the IGM for with lower penetrationredshifts, throughhigher star-formation improvesrates, and older galaxies(longer-duration) bursts of star formation, higher ος, outflows in which the centroid of Lya is shifted redward, and with higher ionizing in the universe."," In general, penetration through the IGM improves for galaxies with lower redshifts, higher star-formation rates, and older (longer-duration) bursts of star formation, higher ${\rm f_{esc}}$, outflows in which the centroid of $\alpha$ is shifted redward, and with higher ionizing backgrounds in the universe."
" Because unabsorbed ionizing photons backgroundseither create Lya in the source or escape to ionize the IGM, but not both, the total emergent Lya decreases for both very high and very low escape fractions."," Because unabsorbed ionizing photons either create $\alpha$ in the source or escape to ionize the IGM, but not both, the total emergent $\alpha$ decreases for both very high and very low escape fractions."
 The fiducial models of Santos(2003) show a broad peak in detectable Lya at fess-0.1—0.8.," The fiducial models of \citet{San03} show a broad peak in detectable $\alpha$ at ${\rm
f_{esc}} \sim 0.1 - 0.8$."
 The strongest constraints on the creation and transmission of Lya come from the cosmic microwave background., The strongest constraints on the creation and transmission of $\alpha$ come from the cosmic microwave background.
" The WMAP observation of the Thompson optical depth to reionization, 7;=0.17+0.4 (Kogutetal.2003),, constrains the parameters in Equation 1.."," The observation of the Thompson optical depth to reionization, $\tau_{e} =
0.17 \pm 0.4$ \citep{Kog03}, constrains the parameters in Equation \ref{eqn:lya}."
" By modeling the reionization process, Cen (2003a,b) and Sokasian et al. ("," By modeling the reionization process, Cen (2003a,b) and Sokasian et al. ("
"2003b) find that a large Thomson optical depth requires most or all of the following at high redshift: a top-heavy IMF of extremely low-metallicity stars, a high star formation efficiency in low-mass halos, fese20.3, a tilt in the matter power spectrum, or an additional source of positiveionizing photons.","2003b) find that a large Thomson optical depth requires most or all of the following at high redshift: a top-heavy IMF of extremely low-metallicity stars, a high star formation efficiency in low-mass halos, $_{\rm esc} \gtrsim 0.3$, a positive tilt in the matter power spectrum, or an additional source of ionizing photons."
" Except for escape fractions extremely close to unity, these conditions are all conducive to observing z>7 galaxies in Lya."," Except for escape fractions extremely close to unity, these conditions are all conducive to observing $z > 7$ galaxies in $\alpha$."
" Although much surrounds the search for Lya in z>7 galaxies, uncertaintywe argue that the distinct possibility of success motivates an attempt with present-day technology."," Although much uncertainty surrounds the search for $\alpha$ in $z > 7$ galaxies, we argue that the distinct possibility of success motivates an attempt with present-day technology."
" As an illustration, we consider two specific scenarios that are both detectable in Lya by design."," As an illustration, we consider two specific scenarios that are both detectable in $\alpha$ by design."
 They differ by a factorof 7-3.4 in the flux of observable Lya photons., They differ by a factorof $\sim$ 3.4 in the flux of observable $\alpha$ photons.
" They are: Both scenarios make assumptions that are relatively favorable for detecting Lya emission, such as the presence of a IMF in high-redshift "," They are: Both scenarios make assumptions that are relatively favorable for detecting $\alpha$ emission, such as the presence of a top-heavy IMF in high-redshift galaxies."
"If regions of intense star top-heavyformation are sufficiently galaxies.chemically enriched by z8 to have a Salpeter IMF from 1—100Mo, Fryaem is reduced by a factor of 4—12, depending on metallicity."," If regions of intense star formation are sufficiently chemically enriched by $z\sim 8$ to have a Salpeter IMF from ${\rm 1-100\ M_{\sun}}$, ${\rm F_{Ly\alpha,em}}$ is reduced by a factor of $4-12$, depending on metallicity."
" With 6.3x105 years of cosmic time available for star formation before z—8, delaying chemical enrichment sufficiently long may require suppressing star formation at early times, as would be the case if the power spectrum were reduced on small scales (e.g.,Yoshidaetal. 2003a,b).."," With $6.3\times 10^8$ years of cosmic time available for star formation before $z = 8$, delaying chemical enrichment sufficiently long may require suppressing star formation at early times, as would be the case if the power spectrum were reduced on small scales \citep[e.g.,][]{Yos03a,Yos03b}. ."
" In addition, inhomogeneous structure formation may lead to pockets of low-metallicity star formation even at late epochs (Scannapieco,Schneider,&Ferrara 2003).."," In addition, inhomogeneous structure formation may lead to pockets of low-metallicity star formation even at late epochs \citep{Sca03}. ."
" Observationally, rest-frame equivalent widths of Lyo as late as z~5.7 support the hypothesis that the IMF may be top-heavy at late epochs (e.g.,Rhoadsetal. 2003).."," Observationally, rest-frame equivalent widths of $\alpha$ as late as $z \sim 5.7$ support the hypothesis that the IMF may be top-heavy at late epochs \citep[e.g.,][]{Rho03}. ."
 More indirect arguments in favor ofmassive and/or early stars come from comparisons of the star formationlow-metallicityhistory ofthe universe to the requirements for reionization, More indirect arguments in favor ofmassive and/or low-metallicity early stars come from comparisons of the star formationhistory ofthe universe to the requirements for reionization
"where W,(0.0) is the window funetion for the j- pixel with the area AQ,.","where $W_p(\theta,\phi)$ is the window function for the $p$ -th pixel with the area $\Delta\Omega_p$ ."
" For the window fuuction IT,(0.0)=1 inside the pixel aud. H,(0.0)—0 outside (Gris et ab 1998). we have from Eq(1)) and Eqt15)): where and The corresponding correlation function (CT98) for the pixclized signal is The disereetuess of the pixelized map determüues the properties of the signal for auy pixels and restricts the precision achieved in any pixclization scheme."," For the window function $W_p(\theta,\phi)=1$ inside the pixel and $W_p(\theta,\phi)=0$ outside (Górrski et $^{6}$ 1998), we have from \ref{eq1}) ) and \ref{eq13}) ): where and The corresponding correlation function (CT98) for the pixelized signal is The discreetness of the pixelized map determines the properties of the signal for any pixels and restricts the precision achieved in any pixelization scheme."
" To estimate this precision we cau use the expansion (CT98) where 5, is the area of the p-th pixel.", To estimate this precision we can use the expansion (CT98) where $S_p$ is the area of the $p$ -th pixel.
 These relatious eeneralize Eq. C». ," These relations generalize Eq. \ref{eq2}) ),"
takiug properties of the window function. iuto account., taking properties of the window function into account.
 The GLESP scheme uses the properties of GaussLegeudre integration in the polar direction while azimuthal pixelizatiou for cach ring is simular to the Igloo scheme aud we ect (see Eq(1))., The GLESP scheme uses the properties of Gauss–Legendre integration in the polar direction while azimuthal pixelization for each ring is similar to the Igloo scheme and we get (see Eq(4)).
" where Ar,EETsy10/2 with c, the p-th GaussLegendre knot aud N the number of pixels in the aziuuthal direction.", where $\Delta x_p=(x_{p+1}-x_{p-1})/2$ with $x_p$ the p-th Gauss–Legendre knot and $N_\phi^p$ the number of pixels in the azimuthal direction.
 This iutegral can be rewritten as follows: where PUTGn) denotes the k-th derivatives at vy., This integral can be rewritten as follows: where ${f^{(k)}}_\ell^m(x_p)$ denotes the k-th derivatives at $x=x_p$ .
" So. for Ae,«Y we get the expausiou of (20)): where where is independent of Αιρ."," So, for $\Delta x_p\ll 1$ we get the expansion of \ref{eq18}) ): where where is independent of $\Delta x_p$."
 For the accuracy of this estimate we eot According to the last inodification of the HIEALDix. an accuracy of the window function reproduction is about 105.," For the accuracy of this estimate we get According to the last modification of the HEALPix, an accuracy of the window function reproduction is about $10^{-3}$."
" To obtain the same accuracy for the (6n). we need to have Using the approximate luk between Legeudre and Bessel facetious for huge (6 (Cradshtevu aud Ryzhil? 2000) we ect: and for Ae,~@/N we have frou Eq.(26)) For mn)exanple. for A—incre we obtain OW(myMG)m23.10 what is a quite reasonable accuracy for (6,447 3000.6000."," To obtain the same accuracy for the $W_p(\ell,m)$, we need to have Using the approximate link between Legendre and Bessel functions for large $\ell$ (Gradshteyn and $^{15}$ 2000) we get: and for $\Delta x_p\sim \pi/N$ we have from \ref{eq24}) ) For example, for $N=2 \ell_{max}$, we obtain $\delta W_p(\ell,m) /
 W_p(\ell,m)\simeq 2.3\times10^{-3}$, what is a quite reasonable accuracy for $\ell_{max}\sim$ 3000–6000."
 The code is developed in two levels of organization., The code is developed in two levels of organization.
 The first one. which uuifies F77 FORTRAN aud € functions. subroutines and wrappers for C routines to be used or FORTRAN calls. cousists of the main procedures ‘signal’ which transforms given values of ἴο a nap. alm which transforms a ap to (iy. Cl2alm which creates a sade of αμ cocticicuts for a given C; and alm2cl which calculaes CV for ay).," The first one, which unifies F77 FORTRAN and C functions, subroutines and wrappers for C routines to be used for FORTRAN calls, consists of the main procedures ' which transforms given values of $a_{\ell m}$ to a map, ' which transforms a map to $a_{\ell m}$, ' which creates a sample of $a_{\ell m}$ coefficients for a given $C_\ell$ and ' which calculates $C_\ell$ for $a_{\ell m}$."
" Procedures or code testing. paraicters coitrol olnogorov-Suiirnov analvsis for Gaussianity of o, and homogencity of phase distribution. aud others. are also inched."," Procedures for code testing, parameters control Kolmogorov-Smirnov analysis for Gaussianity of $a_{\ell m}$ and homogeneity of phase distribution, and others, are also included."
 Operation of hese routines is based ona block of procedures calculating he GaussLegeudre pixclization for a eiven resolution xuanmeter. traustormation of augles to pixel πας a ck.," Operation of these routines is based on a block of procedures calculating the Gauss–Legendre pixelization for a given resolution parameter, transformation of angles to pixel numbers and back."
 The secoud level of the package coutaius the programs which are convenuieut for the 1tilization of the first leve routines., The second level of the package contains the programs which are convenient for the utilization of the first level routines.
 Iu addition to the sraieht use of the already mentioned four nain procedures. they also provide nieans to calculate map patterns generated bv the Joy. 15 and 35» spherical fictions. to compare two sets of τι coefficients.to convert a GLESP map toa IIEALPix map. to convert a IIEALDPix map. or other maps. to a GLESP map Fie.," In addition to the straight use of the already mentioned four main procedures, they also provide means to calculate map patterns generated by the $Y_{20}$, $Y_{21}$ and $Y_{22}$ spherical functions, to compare two sets of $a_{\ell m}$ coefficients,to convert a GLESP map to a HEALPix map, to convert a HEALPix map, or other maps, to a GLESP map Fig."
 6 outlines the CLESP package., \ref{glesp_scheme} outlines the GLESP package.
 The circle defines the zone of the GLESP influence based ou the, The circle defines the zone of the GLESP influence based on the
an improved PSF sampling.,an improved PSF sampling.
 Table 1 stmuuarizes the observations. including imdividual aud total exposure times.," Table \ref{obslog} summarizes the observations, including individual and total exposure times."
 Data reduction lias been performed with IRAE/Pvyraf., Data reduction has been performed with IRAF/Pyraf.
 We combined the individual bias subtracted. flat-fielded nuages with ideutical exposure fines using multidrizzle2002).," We combined the individual bias subtracted, flat-fielded images with identical exposure times using $\texttt{multidrizzle}$."
. This corrects for velocity aberration and geometrical distortion including the 3lth column anomaly based ou the latest distortionL, This corrects for velocity aberration and geometrical distortion including the 34th column anomaly based on the latest distortion.
o For the second epoch observations we applied «2 oversample., For the second epoch observations we applied $\times$ 2 oversampling.
 Astrometry and plotometry were derived from the drizzled images for cach filter aud exposure time settiug using DAOPITOT with a Peuuy2 PSF varvius linearly across the field., Astrometry and photometry were derived from the drizzled images for each filter and exposure time setting using DAOPHOT with a Penny2 PSF varying linearly across the field.
 Near the faint cud. where the photometric uncertainties start fo lucrease. he star list was filled in bv the results derived from he next longer exposure. and uncertainties assessed accordinglv.," Near the faint end, where the photometric uncertainties start to increase, the star list was filled in by the results derived from the next longer exposure, and uncertainties assessed accordingly."
 The fal uuuber of detections iu each baud and epoch for a 56 threshold above the background roise d& listed in Table 1.., The final number of detections in each band and epoch for a $5\sigma$ threshold above the background noise is listed in Table \ref{obslog}.
 Photometric calibration is sed on the zero points provided by Dolphin (2000)., Photometric calibration is based on the zero points provided by Dolphin (2000).
" Photometric correction for charge transfer cficiency ollows the recipe provided byDDolphin’.. aud the astrometric correction is based ou Equation 7 of Wozlurina-Platais et ((2007) with the values for /4. by, aud bs given in Table 2. (note the differcut pixel scale due to 2.2 oversampling for the second epoch)."," Photometric correction for charge transfer efficiency follows the recipe provided by, and the astrometric correction is based on Equation 7 of Kozhurina-Platais et (2007) with the values for $b_1$, $b_2$, and $b_3$ given in Table \ref{param} (note the different pixel scale due to $\times$ 2 oversampling for the second epoch)."
 The combined image has distortious reduced to a level of 0.02 pixel20035)., The combined image has distortions reduced to a level of 0.02 pixel.
". Due to the different oricutation angle of 51"" between the two epochs. we lave to consider the uucertaiutv of the position induced by the residual geometric distortion. Too,=0.017+0.001 pixel."," Due to the different orientation angle of $^\circ$ between the two epochs, we have to consider the uncertainty of the position induced by the residual geometric distortion, $\sigma_{\rm{geo}}=0.017\pm0.001\,\rm{pixel}$ ."
 For sinusoidal pixel phase errors2).. tle second epoch dithering pattern with 0.5 pixel shifts larecly cancels out the pixel phase eror.," For sinusoidal pixel phase errors, the second epoch dithering pattern with 0.5 pixel shifts largely cancels out the pixel phase error."
 As the first epoch was observed in stare mode. the pixel phase error has to be considered.," As the first epoch was observed in stare mode, the pixel phase error has to be considered."
 With a typical amplitude of the sinusoidal Xxel phase error of 0.02 pixel the uncertainty to be included i our analysis amounts to an average residual uncertaiuty of Oyxsph=0.013+0.003 ppixel.," With a typical amplitude of the sinusoidal pixel phase error of $\pm$ 0.02 pixel, the uncertainty to be included in our analysis amounts to an average residual uncertainty of $\sigma_{\rm{pxph}}=0.013\pm0.003$ pixel."
 Sinulatious sed on TiuvTuu PSFEs inelicate that he positional PSF fitting uncertainty results in a ceutroiding error of gpyp=0.013+0.001 pixel., Simulations based on TinyTim PSFs indicate that the positional PSF fitting uncertainty results in a centroiding error of $\sigma_{\rm PSF}=0.013\pm0.001$ pixel.
 The effect ofLST breathing ou pixel scale was determined woincasuringe the separation of wide pairs of stars ou frames obtained during different phases of T's orbit. resulting iu Oprearh=0.009+0.002 pixel.," The effect of breathing on pixel scale was determined by measuring the separation of wide pairs of stars on frames obtained during different phases of 's orbit, resulting in $\sigma_{\rm breath}=0.009\pm0.002$ pixel."
" The combined coutribution of these effects amounts to σι=V↴−TEDTeco|Ophn−ObugOpa,=7qoa)ο2]--aus ο.”n :B", The combined contribution of these effects amounts to $\sigma_{\rm{err}}=\sqrt{\sigma_{\rm{geo}}^2+\sigma_{\rm{pxph}}^2+\sigma_{\rm{PSF}}^2+\sigma_{\rm{breath}}^2}=1.21\pm0.18$ mas.
 observed. proper motion dispersion has to be corrected for σαι to derive the intrinsic velocity dispersion of the cluster members., The observed proper motion dispersion has to be corrected for $\sigma_{\rm err}$ to derive the intrinsic velocity dispersion of the cluster members.
 As the oricutation angles of the two epochs differ, As the orientation angles of the two epochs differ
stellar mass bin in two mass bins.,stellar mass bin in two mass bins.
 We split our sample of 64 ROLES vualaxies into these two bins. divided at 9.2. to give comparable numbers of galaxies in each bin (24 and 40 in the low and high mass bins respectively).," We split our sample of 64 ROLES galaxies into these two bins, divided at $=9.2$ , to give comparable numbers of galaxies in each bin (24 and 40 in the low and high mass bins respectively)."
 Our error bars include a contribution from a reasonable calibration uncertainty(30¢.. the approximate scatter from our comparison with the flux measurements from other To this plot. we add the results of the [OIT]--LD from GDDS 105) for galaxies of higher stellar mass.," Our error bars include a contribution from a reasonable calibration uncertainty, the approximate scatter from our comparison with the flux measurements from other To this plot, we add the results of the -LD from GDDS (J05) for galaxies of higher stellar mass."
 The combination of these two datasets presents a very clear picture of the mass-dependence of [OII]--LD at this redshift., The combination of these two datasets presents a very clear picture of the mass-dependence of -LD at this redshift.
 At a lookback time of around 8 Gyr. the [OII]--LD of the Universe was dominated by high stellar mass galaxies. and a turnover in the [OIT]--LD occurs at )~ 10.0.," At a lookback time of around 8 Gyr, the -LD of the Universe was dominated by high stellar mass galaxies, and a turnover in the -LD occurs at $\sim$ 10.0."
 A simple quadratic fit. for illustrative purposes. would show a peak around )~9.5.," A simple quadratic fit, for illustrative purposes, would show a peak around $\sim$ 9.5."
 Converting the LD to SFRD under our simple model. we measure an integrated SFRD in z-—1 dwarf galaxies (4 < of )9.8pspp=(4.817)10 M. ! Mpe...," Converting the LD to SFRD under our simple model, we measure an integrated SFRD in $\sim$ 1 dwarf galaxies (8.4 $<$ $\le$ 9.8) of $\rho_{SFR} = (4.8\pm1.7) \times 10^{-3}$ $_\odot$ $^{-1}$ $^{-3}$."
" MLFor the first time. we have detectedE the -turnover in he [OII]--LD/SFRD. showing that the contribution of lower mass galaxies. =10""M.. . declines."," For the first time, we have detected the turnover in the -LD/SFRD, showing that the contribution of lower mass galaxies, $\lsim10^{9}$ $_\odot$, declines."
" We construct a local comparison sample from SDSS «data by matching the NYU-VAGC (Blantonetal.2005). sample to flux measurements and stellar masses from the Garching DR4 (Brinchmannetal.2004).. again using the τν, method."," We construct a local comparison sample from SDSS data by matching the NYU-VAGC \citep{Blanton:2005pk} sample to flux measurements and stellar masses from the Garching DR4 \citep{brin04}, again using the $1/V_{max}$ method."
 In order to cleanly sample the [OII|A3727 line. we restrict he redshift range of the sample to 0.082<20.050.," In order to cleanly sample the $\lambda3727$ line, we restrict the redshift range of the sample to $0.032<z\le0.050$."
 We use the same conversion from [OTT]--luminosity to SFR as at ugh redshift. and apply an aperture correction to correct for flux ost outside the fibre using the ratio of the g-band fibre magnitude ο g-band Petrosian magnitude from the imaging data.," We use the same conversion from -luminosity to SFR as at high redshift, and apply an aperture correction to correct for flux lost outside the fibre using the ratio of the $g$ -band fibre magnitude to $g$ -band Petrosian magnitude from the imaging data."
 The local OIIT--determined SFRD is shown as the solid histogram in Fig. 2.., The local -determined SFRD is shown as the solid histogram in Fig. \ref{fig:sfrd}.
 It is clear that the SFRD of low mass galaxies at z~I is comparable to that observed today and that most of the difference is due to a shift in the turnover from high to low mass galaxies with increasing cosmic time., It is clear that the SFRD of low mass galaxies at $\sim$ 1 is comparable to that observed today and that most of the difference is due to a shift in the turnover from high to low mass galaxies with increasing cosmic time.
 It is well known that the dependence of luminosity on the underlying SFR is sensitive to the effects of dust and metallicity (e.g. Jansenetal. 20019).," It is well known that the dependence of luminosity on the underlying SFR is sensitive to the effects of dust and metallicity (e.g., \citealt{jans01}) )."
 The conversion we adopt has been determined empirically from local values (Kennicutt1998). for massive galaxies., The conversion we adopt has been determined empirically from local values \citep{kenn98} for massive galaxies.
 It might be expected that the relationship between luminosity and SFR evolves over the redshift range from z~O-1., It might be expected that the relationship between luminosity and SFR evolves over the redshift range from $\sim$ 0–1.
 Tresseetal.(2002) showed for a modest size sample (30 galaxies) that the ratio of Ha/[OII]. remained constant to within a factor of 2. but with significant uncertainty on the scatter. over this redshift range.," \citet{tres02} showed for a modest size sample (30 galaxies) that the ratio of $\alpha$ remained constant to within a factor of 2, but with significant uncertainty on the scatter, over this redshift range."
 For simplicity. in this Letter we have adopted a single conversion for all our SFR measurements.," For simplicity, in this Letter we have adopted a single conversion for all our SFR measurements."
 It is thus straightforward to convert all our measurements back to -LD and/or adopt a different SFR calibration. if desired.," It is thus straightforward to convert all our measurements back to -LD and/or adopt a different SFR calibration, if desired."
 Given the form of the mass—metallicity relation (e.g. Savaglioetal.2005.. Cowie&Barger 20085) and the metallicity dependence of the luminosity (e.g.. Kewlevetal. 2004). the systematic error in using a single metallicity to estimate [OII]--inferred SFR is that we will the SFR for high stellar mass (thigh metallicity) galaxies andoverestimate the SFR for low mass (low metallicity) objects.," Given the form of the mass–metallicity relation (e.g., \citealt{sava05}, \citealt{cowi08}) ) and the metallicity dependence of the luminosity (e.g., \citealt{kewl04}) ), the systematic error in using a single metallicity to estimate -inferred SFR is that we will the SFR for high stellar mass (high metallicity) galaxies and the SFR for low mass (low metallicity) objects."
 This only strengthens our result that the SFRD declines towards lower mass galaxies at z~ |., This only strengthens our result that the SFRD declines towards lower mass galaxies at $\sim$ 1.
 Our assumption of constant dust-reddening likely leads to a similar systematic error., Our assumption of constant dust-reddening likely leads to a similar systematic error.
 Kewlevetal.(2004) showed from a local galaxy sample that such an assumption leads one to systematically underestimate the SFR at high SFRs and overestimate the SFR at low SFRs., \citet{kewl04} showed from a local galaxy sample that such an assumption leads one to systematically underestimate the SFR at high SFRs and overestimate the SFR at low SFRs.
 This trend also seems to hold for zl galaxies (Adelberger&Steidel20001., This trend also seems to hold for $\gsim$ 1 galaxies \citep{adel00}.
.. Again. the direction of this trend only strengthens our result of a turnover in the SFRD towards lower mass galaxies.," Again, the direction of this trend only strengthens our result of a turnover in the SFRD towards lower mass galaxies."
 The effects of dust obscuration can. be estimated by. for example. looking at the luminosity in the IR. where the emission traces the re-radiated light. due to star- absorbed by dust.," The effects of dust obscuration can be estimated by, for example, looking at the luminosity in the mid-IR, where the emission traces the re-radiated light, due to star-formation, absorbed by dust."
 Conseliceetal.(2007) used 24jim Spüzer observations combined with measurements from DEEP? optical spectroscopy. to attempt to correct [OTT]--inferred SFRs for the effects of dust obscuration.," \citet{cons07} used $\mu$ m observations combined with measurements from DEEP2 optical spectroscopy, to attempt to correct -inferred SFRs for the effects of dust obscuration."
 Their. [OII]2-24;/m SFRs are shown as the open diamonds in Fig. 2..," Their $+$ $\mu$ m SFRs are shown as the open diamonds in Fig. \ref{fig:sfrd},"
 and do indeed suggest that the SFR estimates for high-mass galaxies might be underestimated by [OTT]--only measurements., and do indeed suggest that the SFR estimates for high-mass galaxies might be underestimated by -only measurements.
 Correcting observed SFRs to something approximating SFRs in this way. it can be seen that such corrections can be significant. at least for the high mass galaxies 2:11) sampled by the Conseliceetal.(2007) data.," Correcting observed SFRs to something approximating SFRs in this way, it can be seen that such corrections can be significant, at least for the high mass galaxies $\gsim$ 11) sampled by the \citet{cons07} data."
 Deep 24m data also exist for the CDFS field presented here., Deep $\mu$ m data also exist for the CDFS field presented here.
 Unfortunately. since the surface density of our low mass galaxies is much higher than that of higher mass galaxies. the majority of our sources are limited by confusion and extracting meaningful 24;/m luminosities is likely not possible.," Unfortunately, since the surface density of our low mass galaxies is much higher than that of higher mass galaxies, the majority of our sources are limited by confusion and extracting meaningful $\mu$ m luminosities is likely not possible."
 SFR estimates using other indicators such as flux (and 2-4jim emission. where possible) will be presented for ROLES galaxies in a future paper.," SFR estimates using other indicators such as flux (and $\mu$ m emission, where possible) will be presented for ROLES galaxies in a future paper."
 We note that a potentially small contamination from AGN may exist in our data (although. none of our galaxies is X-ray detected)., We note that a potentially small contamination from AGN may exist in our data (although none of our galaxies is X-ray detected).
 This would also only lower the [OIT]--LD in ROLES., This would also only lower the -LD in ROLES.
 Finally. we overplot the total SFRD (primarily using Ha. but including mass-dependent extinction and metallicity effects) measured locally from SDSS data (Brinchmannetal.2004) in Fig. 2. The Bri," Finally, we overplot the total SFRD (primarily using $\alpha$, but including mass-dependent extinction and metallicity effects) measured locally from SDSS data \citep{brin04} in Fig. \ref{fig:sfrd}."
nchmannetal.(2004) SFR peak occurs at a much higher mass than that of the simple [OTT]--inferred estimate M.)10.5 versus )~ 9.0)., The \citet{brin04} SFR peak occurs at a much higher mass than that of the simple -inferred estimate $\sim10.5$ versus $\sim9.0$ ).
 Given that the SFR traces the Ha SFR. at least on average (Jansenetal.Tresseetal. 2002).. the likely cause of this discrepancy is the use of a single dust extinction factor for the estimate.," Given that the SFR traces the $\alpha$ SFR, at least on average \citep{jans01,tres02}, the likely cause of this discrepancy is the use of a single dust extinction factor for the estimate."
 The individual fits of extinction to each galaxy used by Brinchmannetal.(2004) can vary by several magnitudes., The individual fits of extinction to each galaxy used by \citet{brin04} can vary by several magnitudes.
 Over the relatively small mass range used in ROLES. the dust extinction probably does not vary systematically by a large amount. but in looking at the SFRD over a wider mass range (e.g.. | GDDS). the comparison at low-redshift shows that some care must be taken in interpreting accurately the position of the peak of the SFRD.," Over the relatively small mass range used in ROLES, the dust extinction probably does not vary systematically by a large amount, but in looking at the SFRD over a wider mass range (e.g., $+$ GDDS), the comparison at low-redshift shows that some care must be taken in interpreting accurately the position of the peak of the SFRD."
 Indeed. as we will show in future work. a relatively simple model for mass-dependent extinction. goes most of the way to reconciling the z~O [OTT]--inferred and total SFRDs: however. it is unclear if this is applicable at z—1.," Indeed, as we will show in future work, a relatively simple model for mass-dependent extinction goes most of the way to reconciling the $\sim$ 0 -inferred and total SFRDs; however, it is unclear if this is applicable at $\sim$ 1."
 Given that the conversion from to total SFRDs locally can shift the peak by two orders of magnitude in stellar mass. we are cautious about making quantitative statements regarding the evolution of the peak of the SERD. but our result that the SFRD at z—1 declines towards the lowest mass galaxies 1s robust to any such uncertainties.," Given that the conversion from to total SFRDs locally can shift the peak by two orders of magnitude in stellar mass, we are cautious about making quantitative statements regarding the evolution of the peak of the SFRD, but our result that the SFRD at $\sim$ 1 declines towards the lowest mass galaxies is robust to any such uncertainties."
"To gain further insight iuto the physics of the simulation. we consider the thermal aud chemical properties of the eas, as illustrated iu Figure 5.","To gain further insight into the physics of the simulation, we consider the thermal and chemical properties of the gas, as illustrated in Figure 5."
 We show the free-clectrou abundance. teniperature. and Jeans mass as a function of eas density for every SPID particle.," We show the free-electron abundance, temperature, and Jeans mass as a function of gas density for every SPH particle."
 This πα of presentation coutains an additional dimeusiou of information. concerning the overall timescale of evolution: when the evolution proceeds slowly. particles tend to pile-up in the respective plots. whereas onlv a few particles populate regions im parameter space where the evolution is fast.," This manner of presentation contains an additional dimension of information, concerning the overall timescale of evolution: when the evolution proceeds slowly, particles tend to pile-up in the respective plots, whereas only a few particles populate regions in parameter space where the evolution is fast."
" As can be seen in panel (cl). a gas density of nocPom? and a Jeans mass of AM;~LOCAL, approximately mark the transition between a pliase of slow and fast evolution."," As can be seen in panel (d), a gas density of $n\sim 10^{2}{\rm cm}^{-3}$ and a Jeans mass of $M_{J}\sim 10^{6}M_{\odot}$ approximately mark the transition between a phase of slow and fast evolution."
 Physically. this transition correspouds o the ouset of rnuaway collapse when the clamp mass is ARAM;," Physically, this transition corresponds to the onset of runaway collapse when the clump mass is $M\ga
M_{J}$."
 Iu Figure 6. we show the central eas coufiguratiou for he simulation where the initial spin of the svsteni is ronzero (Run D).," In Figure 6, we show the central gas configuration for the simulation where the initial spin of the system is nonzero (Run B)."
 Iu this case. a binary svstei of clips las formed with a separation of ~10 pe.," In this case, a binary system of clumps has formed with a separation of $\sim
10$ pc."
 These two chuups individually have properties simular to the one chuup formed in Run A. The binary is rapidly diii into he ceuter of the halo by dynamical friction which operates on a timescale fagOre/(CMe) 10* vr (Binney Tremaine 1987)., These two clumps individually have properties similar to the one clump formed in Run A. The binary is rapidly drawn into the center of the halo by dynamical friction which operates on a timescale $t_{\rm df}\sim 0.1 r^{2}v_{c}/(G M_{\rm Cl})\sim 10^{7}$ yr (Binney Tremaine 1987).
 Tere. eis the circular velocity at radius rand Mog is the chip mass.," Here, $v_{c}$ is the circular velocity at radius $r$, and $M_{\rm Cl}$ is the clump mass."
 This timescale is short compared to the Thibble tine at τον10. ἐμz105 vr.," This timescale is short compared to the Hubble time at $z\sim 10$, $t_{H}\simeq
10^{8}$ yr."
 Such asvstem of two compact objects is expected to efiicicuthy radiate eravitational waves that could be detected with the planned (LISA) (Writhe Loeb 20035)., Such a system of two compact objects is expected to efficiently radiate gravitational waves that could be detected with the planned (LISA) (Wyithe Loeb 2003b).
 Massive chuups. either iu isolation or as a binary. are able to form in both runs A aud D. which correspond to different values of initial spin.," Massive clumps, either in isolation or as a binary, are able to form in both runs A and B, which correspond to different values of initial spin."
 The result that such a compact object cau form even in a halo with average iuitial spin (run D). is seenmüuslv at odds with he earlier investigation by Eisenstein Loch (1995). who have argued that only the lowest-spin cosmological rturbatious can harbor seed ΕΠ.," The result that such a compact object can form even in a halo with average initial spin (run B), is seemingly at odds with the earlier investigation by Eisenstein Loeb (1995), who have argued that only the lowest-spin cosmological perturbations can harbor seed BHs."
 The eas that is eventually incorporated iuto the two massive clumps in mu D iust have been able to efficiently lose aueular uonmientun., The gas that is eventually incorporated into the two massive clumps in run B must have been able to efficiently lose angular momentum.
 It is crucial that the evolution iu the A=0.05 case leads to the formation of a binary svstem., It is crucial that the evolution in the $\lambda=0.05$ case leads to the formation of a binary system.
 Tidal orques can then transter mach of the angular moment roni the gas around cach forming chuup to the orbital notion of the system. similar to the case of prescut-dav star formation in a clustered cuvirounent (c.g... Larson 2002: Bate et al.," Tidal torques can then transfer much of the angular momentum from the gas around each forming clump to the orbital motion of the system, similar to the case of present-day star formation in a clustered environment (e.g., Larson 2002; Bate et al."
 2003)., 2003).
 We plan to more fully address he complex physics of augular moneutuni transport iu vinalizig DM halos in future work., We plan to more fully address the complex physics of angular momentum transport in virializing DM halos in future work.
 This efficient mechanism of assembling a large amount of gas into a conipact (1 pe) configuration crucially depends ou the suppresion of star formation in the collapsing halo., This efficient mechanism of assembling a large amount of gas into a compact $\la 1$ pc) configuration crucially depends on the suppresion of star formation in the collapsing halo.
 In Run Ον which otherwise has the sale initial conditions as Run A. we now allow for the formation of hwdrosen molecules.," In Run C, which otherwise has the same initial conditions as Run A, we now allow for the formation of hydrogen molecules."
 Figure 7 Gehich should be compared to Fig., Figure 7 (which should be compared to Fig.
 5 of Run A). shows that the thermal evolution of the eas iu this case proceeds verv differcutly from Ruus A aud D. Here. the gas is able to cool to temperatures Tz200 EK duc to the presence of IL;..," 5 of Run A), shows that the thermal evolution of the gas in this case proceeds very differently from Runs A and B. Here, the gas is able to cool to temperatures $T\ga 200$ K due to the presence of $_{2}$."
 Provided that IIo formation is uot suppressed. the properties of primordial. iietal-free gas are rather similar in svstenis of mass ~105A... and in the smaller halos of mass ~109AF. that are predicted to host the τον first stars at $282030 (see Fie.," Provided that $_2$ formation is not suppressed, the properties of primordial, metal-free gas are rather similar in systems of mass $\sim 10^{8}
M_{\odot}$ and in the smaller halos of mass $\sim 10^{6}M_{\odot}$ that are predicted to host the very first stars at $z\simeq 20 - 30$ (see Fig."
 10 in Broun et al., 10 in Bromm et al.
 2002)., 2002).
 The thermal evolution is clearly reflected in the morphology., The thermal evolution is clearly reflected in the morphology.
 Figure 8 shows the projected σας deusity., Figure 8 shows the projected gas density.
 Due to efficient cooliug. the eas can readilv fragicut and subsequently undergo runaway collapse to form stars already shortlv after turnaround.," Due to efficient cooling, the gas can readily fragment and subsequently undergo runaway collapse to form stars already shortly after turnaround."
 This star formation activity will tend to prevent the assembly of eas in compact clumps in two wavs., This star formation activity will tend to prevent the assembly of gas in compact clumps in two ways.
 First. the Population III stars are expected to be short-lived. with lifetimes of ~3.10° vr. aud they will exert a stroug negative feedback ou their surrouncdiugs upon their death as supernovae (SNe).," First, the Population III stars are expected to be short-lived, with lifetimes of $\sim 3\times 10^{6}$ yr, and they will exert a strong negative feedback on their surroundings upon their death as supernovae (SNe)."
 Second. the SNe will disperse the nucleosvuthetic products from the first generation of stars iuto the remaining eas of the halo.," Second, the SNe will disperse the nucleosynthetic products from the first generation of stars into the remaining gas of the halo."
 Itenceforth. the enriched gas will be able to cool and form stars regardless of whether IT» is preseut or not.," Henceforth, the enriched gas will be able to cool and form stars regardless of whether $_{2}$ is present or not."
 Since the photocdissociation of IT» is now no longer able to lit the eas condensation (see 2). the star formation vate ds expected to merease substantially (ce...Nishi Tashiro 2000: Brom Clarke 2002).," Since the photodissociation of $_{2}$ is now no longer able to limit the gas condensation (see 2), the star formation rate is expected to increase substantially (e.g.,Nishi Tashiro 2000; Bromm Clarke 2002)."
For the low ionization wind. (he intrinsic dipole magnetic fields only confine a small part of the wind.,"For the low ionization wind, the intrinsic dipole magnetic fields only confine a small part of the wind."
 For the high ionization wind in which ions are main composition. the intrinsic magnetic field can severely decrease the mass loss rate.," For the high ionization wind in which ions are main composition, the intrinsic magnetic field can severely decrease the mass loss rate."
 We have developed a multi-fluid model to describe the particle escape from the upper alimosphere of close-in planet., We have developed a multi-fluid model to describe the particle escape from the upper atmosphere of close-in planet.
 The continuity and momentum equations for each component were solved together with an energy equation., The continuity and momentum equations for each component were solved together with an energy equation.
 Detailed micro-physical process. for example. photoionizatlion. recombination and charge exchange. have been included in our mocels.," Detailed micro-physical process, for example, photoionization, recombination and charge exchange, have been included in our models."
 Based on Henvey method. the code can treat svstems of ordinary differential equations wilh one. or more. critical points.," Based on Henyey method, the code can treat systems of ordinary differential equations with one, or more, critical points."
 We calculated two samples of close-in planet. IID 209453b and ILD 139733b.," We calculated two samples of close-in planet, HD 209458b and HD 189733b."
" Our calculations show that the mass loss rates of (wo planets are of the order of magnitude of 10!""—10"" e s! which agree with the observed mass loss rates of LOM—10H ος 1,", Our calculations show that the mass loss rates of two planets are of the order of magnitude of $10^{10}-10^{11}$ g $s^{-1}$ which agree with the observed mass loss rates of $10^{10}-10^{11}$ g $s^{-1}$.
 By detailed test ancl verification. we found that the most important physical process in the atmosphere of hot Jupiter is charge exchange which tightly couples atomic hydrogen with protons.," By detailed test and verification, we found that the most important physical process in the atmosphere of hot Jupiter is charge exchange which tightly couples atomic hydrogen with protons."
 Most of the hydrogen escaping from hot Jupiters is protons. especially in voung star-planet svstem.," Most of the hydrogen escaping from hot Jupiters is protons, especially in young star-planet system."
 Thus. the transit of Lyman-a in WD 189733b could be induced by other unknown physics processes.," Thus, the transit of $\alpha$ in HD 189733b could be induced by other unknown physics processes."
 In the assumption of sinele-fhiid all species must remain collide frequently., In the assumption of single-fluid all species must remain collide frequently.
 Otherwise. the description of sinele-fhiid can not be used.," Otherwise, the description of single-fluid can not be used."
 Our method can be further apply to other case. [or example. (he tenuous wind.," Our method can be further apply to other case, for example, the tenuous wind."
 It is found that decoupling may occur if the mass loss rates are lower than 10? , It is found that decoupling may occur if the mass loss rates are lower than $10^{9}$ 
accretion onto the proto-ueutron star. appareutly provide such conditions (Wanajoetal.2011).,"accretion onto the proto-neutron star, apparently provide such conditions \citep{Wanajo:etal:2011}."
. IToxcever. the 3; obtained under such conditions do not support a strong rprocess. Which successfully reproduces the platinum »eak of r-clements around <=195.," However, the $Y_{e}$ 's obtained under such conditions do not support a strong $r$ -process, which successfully reproduces the platinum peak of $r$ -elements around $A=195$."
 Au alternative explosive scenario as a possible site or the rprocess has been suggested by Joeikunarct where after a supernova explosion a quark phase transition inside the neutrou star can eject very neutrouwrich material from the neutron star surface.," An alternative explosive scenario as a possible site for the $r$ -process has been suggested by \citet{Jaikumar:2007}, where after a supernova explosion a quark--hadron phase transition inside the neutron star can eject very neutron-rich material from the neutron star surface."
 This scenario has been referred to as aquerk-novd.. In he preseut paper. we follow a different approach and investigate core-collapse superunova explosions that are rigeered by a quark-hadrou phase transition during the carly post-bouuce phase (fordetails.seeSagertotal.2009:Fischeretal.2011) and their nucleosvuthesis catures.," This scenario has been referred to as a. In the present paper, we follow a different approach and investigate core-collapse supernova explosions that are triggered by a quark-hadron phase transition during the early post-bounce phase \citep[for details, see][]{Sagert:etal:2009,Fischer:etal:2011}
 and their nucleosynthesis features."
 Under y.such conditious. zones which are initially neutron-rich. can be (promptly) ejected without experienciug the stroug effects of the neutrine flux which conies from the ceutral proto-neutron star.," Under such conditions, zones which are initially neutron-rich, can be (promptly) ejected without experiencing the strong effects of the neutrino flux which comes from the central proto-neutron star."
 Iu the following sections. we discuss the supernova explosion model aud the nuclear plysics inputs utilized for the uucleosvuthesis calculations iu Section 2. the conditions experieuced aud the resulting ejecta iu Section 3. and eive a detailed. analysis of the ejecta composition in Section [. followed by conclusions. also iu comparison to alternative r-process sites. in Section 5.," In the following sections, we discuss the supernova explosion model and the nuclear physics inputs utilized for the nucleosynthesis calculations in Section 2, the conditions experienced and the resulting ejecta in Section 3, and give a detailed analysis of the ejecta composition in Section 4, followed by conclusions, also in comparison to alternative $r$ -process sites, in Section 5."
 We explore uucleosvuthesis in core-collapse superuovae which are triggered by a decoufincment phase trausition during the carly post-bounce phase., We explore nucleosynthesis in core-collapse supernovae which are triggered by a deconfinement phase transition during the early post-bounce phase.
 The explosiou models have been discussed in detail by Fischeretal. (2011)., The explosion models have been discussed in detail by \citet{Fischer:etal:2011}.
. Iu this section. we summarize the main features of the explosion models aud the iuput plysics of the nuclear reaction network utilized for uucleosvuthlesis.," In this section, we summarize the main features of the explosion models and the input physics of the nuclear reaction network utilized for nucleosynthesis."
 A sclfconsistent supernova explosion model. which we adopt to investigate nucleogsvuthesis in the present work. has been carried out by using the AGILE-BOLTZTRAN code.," A self-consistent supernova explosion model, which we adopt to investigate nucleosynthesis in the present work, has been carried out by using the AGILE-BOLTZTRAN code."
 This uunerical code is based on general-relativistic radiation lydrodvuamics d spherical svuuuetry with an adaptive exid and eniplovs a thyec-flavor Boltzimaun neutrino transport (Liebeudórferetal.2001)— with detailed microphysics. e... a realistic nuclear equatiou of state (EOS). weak processes. and (other) nuclear reactions.," This numerical code is based on general-relativistic radiation hydrodynamics in spherical symmetry with an adaptive grid and employs a three-flavor Boltzmann neutrino transport \citep{Liebendoerfer:etal:2004} with detailed microphysics, e.g., a realistic nuclear equation of state (EOS), weak processes, and (other) nuclear reactions."
 For the current study. the staucdard weak. processes were considered as listed in Table 1 of Fischeretal. (2011).," For the current study, the standard weak processes were considered as listed in Table 1 of \citet{Fischer:etal:2011}."
. The barvouic EOS in supernovae needs iu general to be known for the following three iutriusicallv different roelnies: Besides the coutributious from electrons aud positrous ax well as photons. Coulomb corrections to the EOS are added by the method of Times&Aructt(1999).," The baryonic EOS in supernovae needs in general to be known for the following three intrinsically different regimes: Besides the contributions from electrons and positrons as well as photons, Coulomb corrections to the EOS are added by the method of \citet{TimmesArnett:1999}."
 For the current uncleosvuthesis predictions. we select he explosion caleulatious of the 10.8. AL. progenitor nodel where a quarkhiadrou bybrid EOS was used with an carly phase transition to quark matter close to normal miclear matter deusitv (labeledEOS2.secondlineofTable2inFischeretal. 2011).," For the current nucleosynthesis predictions, we select the explosion calculations of the $10.8$ $M_\odot$ progenitor model, where a quark–hadron hybrid EOS was used with an early phase transition to quark matter close to normal nuclear matter density \citep[labeled EOS2, second line of Table 2 in][]{Fischer:etal:2011}."
. We chose this model )ecause it has an explosion cucrev consistent with the expected order of maguitude of 1073 ere (seethesecoudiueofTable3inFischeretal. 2001)., We chose this model because it has an explosion energy consistent with the expected order of magnitude of $10^{51}$ erg \citep[see the second line of Table 3 in][]{Fischer:etal:2011}.
. The masini eravitational mass for this EOS is 1.5026 AL..., The maximum gravitational mass for this EOS is $1.5026$ $M_\odot$.
 Though it is in agreement with the highest precisely known mass of a compact star. the Ππι]κοTavlor pulsar of 1.14 M... recent mass limits for the physical EOS are based ou the willisecond pulsars JL903|0327 with M=1.667ο] AL. (Freiectal.2001). and Ji6112230 with a high nass of M=LOFTEOL AL. (Demorestctal.201001.," Though it is in agreement with the highest precisely known mass of a compact star, the Hulse–Taylor pulsar of $1.44$ $M_\odot$, recent mass limits for the physical EOS are based on the millisecond pulsars J1903+0327 with $M = 1.667 \pm 0.021$ $M_\odot$ \citep{Freire:etal:2011} and J1614–2230 with a high mass of $M = 1.97 \pm 0.04$ $M_\odot$ \citep[][]{Demorest:etal:2010}."
 The inclusion of corrections from thle strong interaction coupling constant can stiffen the quark EOS and lead to üieher maxiuun masses (see.e.g.Schertleretal.2000:Weisscnboruetal. 2011).," The inclusion of corrections from the strong interaction coupling constant can stiffen the quark EOS and lead to higher maximum masses \citep[see, e.g.,][]{Schertler00,Alford:etal:2005,Sagert:etal:2010,Sagert:etal:2011,Weissenborn:etal:2011}."
. The physical conditions for a possible quark-radron phase transition iu the proto-ueutron/ star— are Helly uncertain., The physical conditions for a possible quark-hadron phase transition in the proto-neutron star are highly uncertain.
 Therefore. Fischerotal.(2011) also constructed a quark-hadron hybrid EOS which includes corrections frou the strong interaction coupling constant.," Therefore, \citet{Fischer:etal:2011} also constructed a quark-hadron hybrid EOS which includes corrections from the strong interaction coupling constant."
" They showed that such an EOS. with a niaxinmuni mass of 1.67 AL. aud a phase transition to quark iatter close to nuclear saturation deusity (sceEOS3intable2and3Fischeretal.2011).. leads to a qualitatively similar explosion scenario in spherical svnunetry as obtained in the explored models with several EOSs which have lower ιαππα masses,"," They showed that such an EOS, with a maximum mass of 1.67 $_\odot$ and a phase transition to quark matter close to nuclear saturation density \citep[see EOS3 in table~2 and 3 in][]{Fischer:etal:2011}, leads to a qualitatively similar explosion scenario in spherical symmetry as obtained in the explored models with several EOSs which have lower maximum masses."
 Note. the different post-bounce times for the ouset of cecoufinement. and hence for the onset of the explosion. lead to differeut proto-neutron star structures.," Note, the different post-bounce times for the onset of deconfinement, and hence for the onset of the explosion, lead to different proto-neutron star structures."
 Tn general. a delaved plase transition at higher deusities translates to a longer mass accretion phase (for the same progenitor).," In general, a delayed phase transition at higher densities translates to a longer mass accretion phase (for the same progenitor)."
" It results iu a iore massive protoucutrou star with a steeper density eradicut at its surface aud a lower Y, at the ouset of the explosion triggered by the phase transition.", It results in a more massive protoneutron star with a steeper density gradient at its surface and a lower $Y_e$ at the onset of the explosion triggered by the phase transition.
 An earlier phase transition aud less, An earlier phase transition and less
used for the continuum underlying the X-ray light curve of that sample are thesame.,used for the continuum underlying the X-ray light curve of that sample are the.
 If we classify the light curves with flares (FLC) on the basis of their best-fitting function as in Sect., If we classify the light curves with flares (FLC) on the basis of their best-fitting function as in Sect.
" ?? we find that the sample of Marguttietal.(2011) contains a majority of Type II light curves (72%), 14% of Type 0, and 7% of Types Ia and Ib."," \ref{class} we find that the sample of \citet{2011MNRAS.410.1064M} contains a majority of Type II light curves $72\%$ ), $14\%$ of Type 0, and $7\%$ of Types Ia and Ib."
" This repartition is different from what we found in our sample (34% of Type 0, 31% of Type Ia, 11% of Type Ib, and only 23% of Type II, see also Marguttietal. 2010))."," This repartition is different from what we found in our sample $34\%$ of Type 0, $31\%$ of Type Ia, $11\%$ of Type Ib, and only $23\%$ of Type II, see also \citealt{giantflares10}) )."
" We produced a population of ""fake"" flares on the basis of the relations found in Chincarinietal.(2010) (w=0.2tj--10; tdee~2trises Loko Eftares~10°! erg) with 40 s X{νι<500 s, and we DEsuperimposed these flares to the light curves of our golden sample."," We produced a population of “fake” flares on the basis of the relations found in \citet{chinca10} $w=0.2\,t_{pk}+10$; $t_{dec}\sim2\,t_{rise}$; $L_{pk}\propto t_{pk}^{-2.7}$; $E_{flares}\sim10^{51}$ erg) with $40$ s $\leqslant t_{pk}\leqslant 500$ s, and we superimposed these flares to the light curves of our golden sample."
 We find that a “standard” flaring activity would have been detected above lo in 80% of the golden sample light curves., We find that a “standard” flaring activity would have been detected above $1 \sigma$ in $80\%$ of the golden sample light curves.
" In particular, these flares could have been detected in all the five Type Ia light curves for t=100 s. Moreover, if we compute the average luminosity of the light curves of the golden sample, we find that it has a shallower decay ((L)ο. f?) but is even dimmer than the average continuum obtained in Marguttietal.(2011) for t€200 s. This excludes a luminosity bias preventing detection of early time flares in the light curves of our golden sample."," In particular, these flares could have been detected in all the five Type Ia light curves for $t \gtrsim 100$ s. Moreover, if we compute the average luminosity of the light curves of the golden sample, we find that it has a shallower decay $\left\langle L\right\rangle\propto t^{-2}$ ) but is even dimmer than the average continuum obtained in \citet{2011MNRAS.410.1064M} for $t \lesssim 200$ s. This excludes a luminosity bias preventing detection of early time flares in the light curves of our golden sample."
 In Fig., In Fig.
 4 we compare the distribution of the steep decay power-law index of Type II and Ib light curves of the present sample (LC) with the corresponding value found for the FLC., \ref{alpha} we compare the distribution of the steep decay power-law index of Type II and Ib light curves of the present sample (LC) with the corresponding value found for the FLC.
 We did not include those cases in which the prompt emission still observed in the X-rays (as GRB060510B)., We did not include those cases in which the prompt emission still observed in the X-rays (as GRB060510B).
" The FLC power-law index distribution has a median (a?LC)=2.7 (g*= 1.3), while the LC sample has (a)-32 (c.€= 1.2)."," The FLC power-law index distribution has a median $\left\langle \alpha^{FLC}\right\rangle=2.7$ $\sigma^{LC}=1.3$ ), while the LC sample has $\left\langle \alpha^{FLC}\right\rangle=3.2$ $\sigma^{LC}=1.2$ )."
 The K-S test comparing the temporal indices a! and a! gives a probability of 9% that they belong to the same population., The K-S test comparing the temporal indices $\alpha^{LC}$ and $\alpha^{FLC}$ gives a probability of $9\%$ that they belong to the same population.
 Marguttietal.(2011) show that light curves with multiple flares tend to be shallower than those with a single flare., \citet{2011MNRAS.410.1064M} show that light curves with multiple flares tend to be shallower than those with a single flare.
 We then select the FLC with a single flare that are expected to be more similar to LCs and compare their power-law indices with the LC distribution., We then select the FLC with a single flare that are expected to be more similar to LCs and compare their power-law indices with the LC distribution.
" The two distributions are more similar, with a median of the single FLC distribution (72€)=3.0 (fC= 1.0)."," The two distributions are more similar, with a median of the single FLC distribution $\left\langle \alpha_{sing}^{FLC}\right\rangle=3.0$ $\sigma_{sing}^{FLC}=1.0$ )."
 The K-S test comparing the temporal indices a/© and results in an increasing probability that they belong to the Onesame population (27%)., The K-S test comparing the temporal indices $\alpha^{LC}$ and $\alpha_{sing}^{FLC}$ results in an increasing probability that they belong to the same population $27\%$ ).
 We aim at investigating now the late-time behaviour of the entire sample., We aim at investigating now the late-time behaviour of the entire sample.
 For this purpose we computed the average light curve for the entire LC sample., For this purpose we computed the average light curve for the entire LC sample.
" After 1000 s we observe that (L)«Γ12, a trend similar to the detection threshold found for the average flare light curve in Marguttietal.(2011).."," After $1000$ s we observe that $\left\langle L\right\rangle\propto t^{-1.2}$, a trend similar to the detection threshold found for the average flare light curve in \citet{2011MNRAS.410.1064M}."
" If we compare the average luminosity of the LC sample with the one of the continuum underlying the late-time flares (LFLC) of Bernardinietal.(2011),, we find that the temporal behaviour is very similar but the present sample is brighter, with ΤάΕΙΟν2096."," If we compare the average luminosity of the LC sample with the one of the continuum underlying the late-time flares (LFLC) of \citet{2011A&A...526A..27B}, we find that the temporal behaviour is very similar but the present sample is brighter, with $L^{LFLC}/L^{LC}\sim 20\%$."
 We turn now to the energy output of the X-ray light curves with flares., We turn now to the energy output of the X-ray light curves with flares.
" We restrict our analysis to a (FGS) of 17 light curves, defined with the same criterion as was introduced in Sect. ??,,"," We restrict our analysis to a (FGS) of $17$ light curves, defined with the same criterion as was introduced in Sect. \ref{sample},"
 and we calculate the energy of the underlying continuum in the common rest frame energy band (1.9—11.2 keV) as in Sect. ??.., and we calculate the energy of the underlying continuum in the common rest frame energy band $1.9-11.2$ keV) as in Sect. \ref{energy}.
 We find that the distributions of both the X-ray energy and the isotropic prompt emission energy are similar for the two samples., We find that the distributions of both the X-ray energy and the isotropic prompt emission energy are similar for the two samples.
" A correlation between the E, and Ej;,, is present (p=0.70, see Fig. 3)),"," A correlation between the $E_x$ and $E_{iso}$ is present $\rho=0.70$, see Fig. \ref{etype}) ),"
 as well as between the total fluence in the XRT energy band (0.3—10 keV) and the prompt emission fluence detected by BAT (15—150 keV) for the complete sample of long GRBs with and without flares from the XRT Catalogue by Margutti et al. (, as well as between the total fluence in the XRT energy band $0.3-10$ keV) and the prompt emission fluence detected by BAT $15-150$ keV) for the complete sample of long GRBs with and without flares from the XRT Catalogue by Margutti et al. (
in preparation) with a complete light curve in the sense described in the previous section (p= 0.51).,in preparation) with a complete light curve in the sense described in the previous section $\rho=0.51$ ).
" The correlation between Eyj,,, and Ejso is enhanced including the FGS data (see Fig. 3.", The correlation between $E_{plateau}$ and $E_{iso}$ is enhanced including the FGS data (see Fig. \ref{etype}.
.e)., .e).
 The results of previous section are confirmed by adding the light curves with flares also in the fluence-fluence plot for the plateau phase (p— 0.50) and the shallow decay (p= 0.63)., The results of previous section are confirmed by adding the light curves with flares also in the fluence-fluence plot for the plateau phase $\rho=0.50$ ) and the shallow decay $\rho=0.63$ ).
 The other fluence-fluence relations portrayed in Fig., The other fluence-fluence relations portrayed in Fig.
3 are not improved significantly with the exception of the Type 0 versus prompt emission (p= 0.50)., \ref{etype} are not improved significantly with the exception of the Type 0 versus prompt emission $\rho=0.50$ ).
 The analysis of the rest-frame X-ray light curves of our sample shows that they can be grouped into four morphological types., The analysis of the rest-frame X-ray light curves of our sample shows that they can be grouped into four morphological types.
" We found that the steep decay of Types Ib and II can be grouped together, as well, as the shallow decays in Types Ia, Ib, and IL and normal decays in Types 0, Ia, and Π."," We found that the steep decay of Types Ib and II can be grouped together, as well, as the shallow decays in Types Ia, Ib, and II, and normal decays in Types 0, Ia, and II."
" Spectral evolution is associated with the steep decay, while no spectral evolution is found during shallow and normal decay."," Spectral evolution is associated with the steep decay, while no spectral evolution is found during shallow and normal decay."
 We demonstrated that, We demonstrated that
Eq. (,Eq. (
12) for O44 eives This equation is consistent with the analytical results obtained by Desai [7]..,12) for $O_{13}$ gives This equation is consistent with the analytical results obtained by Desai \cite{Desai}.
 Eq. (, Eq. (
14) can be used to obtain the relation as shown in our earlier work |11]..,14) can be used to obtain the relation as shown in our earlier work \cite{mee paper}.
 Using Eq. (, Using Eq. (
23) and neglecting the higher order terms in 543. Eq. (,"23) and neglecting the higher order terms in $s_{13}$, Eq. ("
22) can be written as ,"22) can be written as For the tri-bi-maximal value \cite{TBM} $s_{12}^2=\frac{1}{3}$, the trigonometric factor in Eq. ("
Substituting the present best [it values of the oscillation parameters in Eq. (,24) is given by Substituting the present best fit values of the oscillation parameters in Eq. (
"24) . Am},=?4x10* and sq»=0.3 131). we find that 944=7.3° which is consistent with the earlier estimates |4.5]. for this mass matrix texture.","24) $\Delta m^2_{12}=7.9\times10^{-5}$ , $\Delta m^2_{13}=2.4\times10^{-3}$, and $s_{12}=0.3$ \cite{fogli}) ), we find that $\theta_{13}= 7.3^o$ which is consistent with the earlier estimates \cite{a1a2 paper,2t paper}
 for this mass matrix texture."
 Using the above leading order estimates of the elements 4A... D. C and D. the neutrino mass malrix given in Eq. (," Using the above leading order estimates of the elements $A$, $B$, $C$ and $D$, the neutrino mass matrix given in Eq. ("
1) can be written as where (he parameter e is given by for maximal alimospheric mixing.,1) can be written as where the parameter $\epsilon$ is given by for maximal atmospheric mixing.
 For the present best fit values of the neutrino oscillation parameters. e~0.36.," For the present best fit values of the neutrino oscillation parameters, $\epsilon\sim 0.36$."
 So. the hierarchical structure of neutrino mass matrix should be as follows: The phenomenological neutrino mass matrix obtained here has to be confronted with the neulrino mass matrices given by various neutrino mass models.," So, the hierarchical structure of neutrino mass matrix should be as follows: The phenomenological neutrino mass matrix obtained here has to be confronted with the neutrino mass matrices given by various neutrino mass models."
 As an example. we confront," As an example, we confront"
have shown delays iu the onset of the LAT component with respect to the MeV cimission.,have shown delays in the onset of the LAT component with respect to the MeV emission.
 Understaudiug the origin of these behaviors will provide clues for a better understaudius of the prompt emission of GRBs., Understanding the origin of these behaviors will provide clues for a better understanding of the prompt emission of GRBs.
in the elliptical galaxies. such as NGC 4261. because they generally do not exhibit substantial rotation. although even initially slow rotation would lead to substantial rotation il the dust settled into the central kiloparsec from Large scales.,"in the elliptical galaxies, such as NGC 4261, because they generally do not exhibit substantial rotation, although even initially slow rotation would lead to substantial rotation if the dust settled into the central kiloparsec from large scales."
 Mathews&Drighenü(2003) model the evolution of the clusty gas originating in stellar winds and [found cooling dust-rich. gas could produce (the observed. cireumnmiclear dust. and in particular the observed. clumping.," \citet{mathews03} model the evolution of the dusty gas originating in stellar winds and found cooling dust-rich gas could produce the observed circumnuclear dust, and in particular the observed clumping."
 llowever. even if dust cooling explains the observed clumpiness. stellar mass loss still can nol explain why cireumnuclear dust is only present in ~50% of all early-type galaxies.," However, even if dust cooling explains the observed clumpiness, stellar mass loss still can not explain why circumnuclear dust is only present in $\sim 50$ of all early-type galaxies."
 While dust destruction via collisions with hot gas will occur in on order 105 vears 2003).. mass loss will continually supply new dust (o the interstellar medium.," While dust destruction via collisions with hot gas will occur in on order $10^8$ years \citep[e.g.,][]{mathews03}, mass loss will continually supply new dust to the interstellar medium."
 This difficulty with internal dust production is also rellected in (he absence of a correlation between [ar-intrarecl and visible-waveleneth huminosities in earlv-0(vpe galaxies. which also suggests (he dust may have an external origin (Temietal.2004).," This difficulty with internal dust production is also reflected in the absence of a correlation between far-infrared and visible-wavelength luminosities in early-type galaxies, which also suggests the dust may have an external origin \citep{temi04}."
. Kinematic observations of neutral and ionized gas in early-(vpe galaxies have led many authors to conclude that this material has an external origin because the gas kinematics diller from the stellar kinematics (e.g.Bertolaetal.1984.1992:Sarzi2006:Morganti2006.andreferences therein)..," Kinematic observations of neutral and ionized gas in early-type galaxies have led many authors to conclude that this material has an external origin because the gas kinematics differ from the stellar kinematics \citep[e.g.][and 
references therein]{bertola84,bertola92,sarzi06,morganti06}."
 The circumnuclear morphology of the dust also provides some additional constraints on the origin of the dust., The circumnuclear morphology of the dust also provides some additional constraints on the origin of the dust.
 In particular. (he chaotic appearance suggests an external origi because internallv-created dust might be too uniformly distributed. particularly in lenticulars. if the cooling time discussed above is not shorter than the dvnamical (ime.," In particular, the chaotic appearance suggests an external origin because internally-created dust might be too uniformly distributed, particularly in lenticulars, if the cooling time discussed above is not shorter than the dynamical time."
 llowever. the circumnuclear. dust. observed in late-tvpe galaxies exhibits a similar range in morphologies and that material is generally not held to have an external origin," However, the circumnuclear dust observed in late-type galaxies exhibits a similar range in morphologies and that material is generally not held to have an external origin."
 The morphology of the dust can also provide some constraints on the lifetime of the dust., The morphology of the dust can also provide some constraints on the lifetime of the dust.
 discuss how the different. cust morphologies can be viewed as a settling sequence. where dust that is not in a disk must settle toward the center of the galaxy. in a dvnanmical time. or a lew 10* vears (seealsoTranοἱal.2001: 2005).. while coherent. tightly wound dust disks must be at least a few dynamical times old.," \citet{lauer05} discuss how the different dust morphologies can be viewed as a ""settling sequence,"" where dust that is not in a disk must settle toward the center of the galaxy in a dynamical time, or a few $10^7$ years \citep[see also][]{tran01,kleijn05}, while coherent, tightly wound dust disks must be at least a few dynamical times old."
 The age of such a dust disk would then be set by the dust destruction timescale mentioned above. or eI0? vears.," The age of such a dust disk would then be set by the dust destruction timescale mentioned above, or $\sim 10^8$ years."
 External dust input and a finite cust lifetime can plausibly explain why cust is not found in the centers of all earlv-tvpe galaxies., External dust input and a finite dust lifetime can plausibly explain why dust is not found in the centers of all early-type galaxies.
 More broadly. the relative number of galaxies with chaotic dust lanes. some nuclear dust spirals. and verv coherent dust disks could provide constraints on (he settling time of the dust if there is a natural progression Coward more ordered dust as the dust settles toward the galaxy nucleus or is destroved.," More broadly, the relative number of galaxies with chaotic dust lanes, some nuclear dust spirals, and very coherent dust disks could provide constraints on the settling time of the dust if there is a natural progression toward more ordered dust as the dust settles toward the galaxy nucleus or is destroyed."
 If this model is correct. then the dust in inactive galaxies should be most chaotic or unsettled.," If this model is correct, then the dust in inactive galaxies should be most chaotic or unsettled."
 Our data are consistent will this interpretation. with the exception of the tightly. wound dust spiral in NGC 4526. although there are only a small inunber of tightly wound dust spirals in our full earlv-tvpe galaxy saniple.," Our data are consistent with this interpretation, with the exception of the tightly wound dust spiral in NGC 4526, although there are only a small number of tightly wound dust spirals in our full early-type galaxy sample."
 In fact. the relative scarcity of the tightly wound dust spirals presents a problem because (he estimates," In fact, the relative scarcity of the tightly wound dust spirals presents a problem because the estimates"
the individual systems being members of a homogeneous sample or not.,the individual systems being members of a homogeneous sample or not.
 Let /7;(0;|£) and PotCM/L)£) be the probabilities of the data given the parameters for galaxy 7 if it is a member of a homogeneous group based on (he measured. “good uncertsinties. and Pyjto;|£) and Py;(CU/L);£) be the probabilities if it is not ancl we should be using the “bac” uncertainties based on the svstematic error estimates clerived from our first method.," Let $P_{Gi}(\sigma_i|\bfxi)$ and $P_{Gi}((M/L)_i|\bfxi)$ be the probabilities of the data given the parameters for galaxy $i$ if it is a member of a homogeneous group based on the measured, “good” uncertainties, and $P_{Bi}(\sigma_i|\bfxi)$ and $P_{Bi}((M/L)_i|\bfxi)$ be the probabilities if it is not and we should be using the “bad” uncertainties based on the systematic error estimates derived from our first method."
 The Press(1997) provides estimates of (he relative likelihoods describing either the full sample or the individual svstems by (he “good” or bad data model., The \citet{bo97} provides estimates of the relative likelihoods describing either the full sample or the individual systems by the “good” or “bad” data model.
 If. we want the Bavesian probability distribution for the parameters € properly weighted over all possible eroup membership combinations. we find that where and where P(£) sets the prior probability distributions for the parameters.," If we want the Bayesian probability distribution for the parameters $\bfxi$ properly weighted over all possible group membership combinations, we find that where and where $P(\mathbf{\xi})$ sets the prior probability distributions for the parameters."
" We assume a uniform priors for p, aud pp.", We assume a uniform priors for $p_\sigma$ and $p_L$.
 We obtain the probability distribution for anv parameter by marginalizing Eqn. (7)), We obtain the probability distribution for any parameter by marginalizing Eqn. \ref{eqn:p1}) )
 over all other variables ancl (hen normalizing the total probability to unity., over all other variables and then normalizing the total probability to unity.
 We can also estimate the probability that the sample is homogeneous in eitlier its structural or evolutionary properties as and the probability (hat a particular galaxy is in the dynamically homogeneous class is where A similar set of expressions oOgives the probability that the egalaxy is in the set of egalaxies with a homogeneous evolutionary history., We can also estimate the probability that the sample is homogeneous in either its structural or evolutionary properties as and the probability that a particular galaxy is in the dynamically homogeneous class is where A similar set of expressions gives the probability that the galaxy is in the set of galaxies with a homogeneous evolutionary history.
is not diagonal.,is not diagonal.
 Iu gencral. Cz is nof evel ciagonalizable.," In general, $\Cb_{\zhb}$ is not even diagonalizable."
 For example. this happens when the oxiodoerani of a time series contaiuiug Af data is computed on JN frequencies with Α>AM (a typical situation in practical applications).," For example, this happens when the periodogram of a time series containing $M$ data is computed on $N$ frequencies with $N > M$ (a typical situation in practical applications)."
 This nuplies tha he entries of Z cannot be made mutually uncorrelated., This implies that the entries of $\zhb$ cannot be made mutually uncorrelated.
 Obviously. the same holds for the eutries of p as eiven w Eq. (15)).," Obviously, the same holds for the entries of $\phb$ as given by Eq. \ref{eq:period3}) )."
 As a consequence. although a iuuber NV of yequencies ave considered in p. a most only AY/2 of them are statistically indepeudeut?.," As a consequence, although a number $N$ of frequencies are considered in $\phb$, at most only $M/2$ of them are statistically ."
". Particularly troublesome is hat. for a given frequency A. Riv]. aud Z[6,] ty, ane oy, pe) are also correlated."," Particularly troublesome is that, for a given frequency $k$, $\Rmatc[\xhu_k]$, and $\Imatc[\xhu_k]$ $\zh_k$ and $\zh_{N_{\dag}+k}$ ) are also correlated."
 This makes it difficult to fix the statistica characteristics of pj., This makes it difficult to fix the statistical characteristics of $\ph_k$.
" Iu this respect. two choices are yossible,"," In this respect, two choices are possible."
 The first consists in the determination. for cach frequency. of he PDF of py.," The first consists in the determination, for each frequency, of the PDF of $\ph_k$ ."
" Actually. this is a rather involved approach. because Z4/ and Ty.)lip have variance (6πλ and (Cz),piv,ee respectively. aud covariauce (τν,κ."," Actually, this is a rather involved approach, because $\zh_k$ and $\zh_{N_{\dag} + k}$ have variance $(C_{\zhb})_{kk}$ and $(C_{\zhb})_{N_{\dag} + k, N_{\dag} +
 k}$, respectively, and covariance $(C_{\zhb})_{k, N_{\dag} +
 k}$."
" Therefore. once 2; and Ty,jj ave normalized to unif variance. each p; is given by the sun of two correlated d random quantities"," Therefore, once $\zh_k$ and $\zh_{N_{\dag} + k}$ are normalized to unit variance, each $\ph_k$ is given by the sum of two correlated $\chi_1^2$ random quantities."
 Although available in analvtical form (Simon2006).. the resulting PDF is rather conrplex and hence difficult to handle (foranalterna-tiveapproach.seeReeseu 2007).," Although available in analytical form \citep{sim06}, the resulting PDF is rather complex and hence difficult to handle \citep[for an alternative approach, see][]{ree07}."
". Moreover. there is the additional problems that Z;, changes with A."," Moreover, there is the additional problem that $L_{\ph_k}$ changes with $k$."
" À simpler alternative is the use of two wncorrelated aud umit-variauce random quantities. Pj and My,τν. obtained through the transformation Tere,e οopi—[6rAAON,TN ae2ee−−~N;TN M.1 D>. Inia diagonal matrix whose eutries are given by the reciprocal of the square root of the non-zero eigenvalues of the covariuice nnmtrix zero otherwise. aud U, is an orthogonal matrix that contains the correspondingcigeuvectors’."," A simpler alternative is the use of two uncorrelated and unit-variance random quantities, $\uph_k$ and $\uph_{N_{\dag} + k}$, obtained through the transformation Here, $\uphb^T = [\uph_k, \uph_{N_{\dag} + k}]$, $\zhb_\star^T =
[\zh_k, \zh_{N_{\dag} + k}]$, $\Sigmab_{\star}^{-1/2}$ is a diagonal matrix whose entries are given by the reciprocal of the square root of the non-zero eigenvalues of the covariance matrix zero otherwise, and $\Ub_\star$ is an orthogonal matrix that contains the corresponding."
". Indeed. if the periodogran is defined as then cach p, is given by the suu of two independent. unit-variance. Caussian random quantities."," Indeed, if the periodogram is defined as then each $\ph_k$ is given by the sum of two independent, unit-variance, Gaussian random quantities."
 As a consequence. the corresponding PDF is indepencently of ke a 3 whose (CDE) is the exponential function.," As a consequence, the corresponding PDF is, independently of $k$ , a $\chi^2_2$ whose (CDF) is the exponential function."
 This permits determining the statistical significance of pj for aspecified frequeucy k., This permits determining the statistical significance of $\ph_k$ for a frequency $k$ .
 Things become muore complex if Wy frequeucies are inspected when lookiug for a peak., Things become more complex if $N_f$ frequencies are inspected when looking for a peak.
 Indeed. also after the operation (26)]. it happens that ρεz0 for &zf. ic. the frequencies of the periocogram remain mutually correlated.," Indeed, also after the operation \ref{eq:norm}) ), it happens that ${\rm E}[\ph_k \ph_l] \neq 0$ for $k \neq l$, i.e., the frequencies of the periodogram remain mutually correlated."
 This is an unavoidable problem., This is an unavoidable problem.
 Because of it. Ny does not correspoud to the uuuber offrequencies. so the (20)) cannot be coniputed.," Because of it, $N_f$ does not correspond to the number of, so the \ref{eq:false}) ) cannot be computed."
" Iowever. since £5, is the sae for all the frequencies. al upper iuit can be fixed for Zg, by setting Αγ[A//2|]."," However, since $L_{\ph_k}$ is the same for all the frequencies, an upper limit can be fixed for $L_{{\rm Fa}}$ by setting $N_f = \lceil
M/2 \rceil$."
 The periodogram obtained by nemus of Eq. (28)), The periodogram obtained by means of Eq. \ref{eq:step1}) )
 corresponds to the originalLonib-Seeryle periodogran., corresponds to the original periodogram.
 Since the transformation (17)) docs not depeud on the characteristics of the signal sampling. the stratceey of felowing in the case that a is the realization of (110 jecessarilv stationary) colored uoise is simply the one in Sec.," Since the transformation \ref{eq:transf}) ) does not depend on the characteristics of the signal sampling, the strategy of following in the case that $\xb$ is the realization of (not necessarily stationary) colored noise is simply the one in Sec."
 2. ttraustormation of a to an array gy with uncorrelated entries., \ref{sec:formalization} transformation of $\xb$ to an array $\yb$ with uncorrelated entries.
 After that. theLomb-Scargle periodogranmià can be computed.," After that, the periodogram can be computed."
 It is worth noticing that this simple result has been possible thanks to a formulation of the problem iu the time domain aud the use of the natrix notation., It is worth noticing that this simple result has been possible thanks to a formulation of the problem in the time domain and the use of the matrix notation.
 The same results could have Όσοι obtained by following the popular approach of working iu the harmouwic domain but at the price of a much more cifficul derivation., The same results could have been obtained by following the popular approach of working in the harmonic domain but at the price of a much more difficult derivation.
 To illustrate the usefuluess and the simplicity of the proposed formalisin in handling different situations from the classical oucs. we show two examples in this section.," To illustrate the usefulness and the simplicity of the proposed formalism in handling different situations from the classical ones, we show two examples in this section."
 The first consists of a periodoerani of a nmeau-subtracted time series., The first consists of a periodogram of a mean-subtracted time series.
 The evaluation of he reliability of a peak in the periodograim of a signal à requires that (uuder themidi hiypothesis $= 0n) n be the realization of a zero-nean boise process., The evaluation of the reliability of a peak in the periodogram of a signal $\xb$ requires that (under the hypothesis $\xb = \nb$ ) $\nb$ be the realization of a zero-mean noise process.
 Iu most experimental situations. this condition is not fulfilled aud oue works with a centered muuean-subtracted) version x of ax.," In most experimental situations, this condition is not fulfilled and one works with a centered mean-subtracted) version $\chib$ of $\xb$."
 This. however. introduces some (often neglected) problems.," This, however, introduces some (often neglected) problems."
 The case where a is the realization of a discrete white noise process with: variance: oF9 has been cousicered: several times: in: the literature., The case where $\xb$ is the realization of a discrete white noise process with variance $\sigma^2_{\xb}$ has been considered several times in the literature.
 An example is the paper by Zechuecister&Ister(2009) where a rather elaborate solution is presented., An example is the paper by \citet{zec09} where a rather elaborate solution is presented.
 With the approach proposedhere. a simpler solution can be obtained if one takesinto. consideratiou that the subtraction of the mean from a forces a spurious correlation among the eutries of xiu sucha way that the covariauice uatrix Cy=E|xx?| is given by where AY is umber of entries of a and Tan ArxALmatrix with every eutry equal to unity;," With the approach proposedhere, a simpler solution can be obtained if one takesinto consideration that the subtraction of the mean from $\xb$ forces a spurious correlation among the entries of $\chib$in sucha way that the covariance matrix $\Cb_{\chib}=E[ \chib \chib^T ]$ is given by where $M$ is number of entries of $\xb$ and $\large{\oneb}$ an $M \times M$matrix with every entry equal to unity."
 Since this matrix is, Since this matrix is
The physical properties of laminar helium deflagrations. which are primarily cletermined by a balauce between nuclear euergy generation and the treisport of internal energy. have been evaluated by Tinuues (1999) for a large grid of upstream cdeusities aud temperatures.,"The physical properties of laminar helium deflagrations, which are primarily determined by a balance between nuclear energy generation and the transport of internal energy, have been evaluated by Timmes (1999) for a large grid of upstream densities and temperatures."
 We will use tle results ol this survey for values of the laminar deflagration speeds δι. widths 9. and density. contrasts Ap/p between the unburned fuel aud its ash.," We will use the results of this survey for values of the laminar deflagration speeds $S_{\rm L}$, widths $\delta$, and density contrasts $\Delta \rho/\rho$ between the unburned fuel and its ash."
 The density of the ash is smaller than the density of the unburued fuel because the density declines behind a subsonic flame front., The density of the ash is smaller than the density of the unburned fuel because the density declines behind a subsonic flame front.
 There are several plausible definitions for the width of the deflagration., There are several plausible definitions for the width of the deflagration.
 We consider tluee pragmatic definitious. each of which can be measured by resolved. calculatious of laminar helium deflagrations.," We consider three pragmatic definitions, each of which can be measured by resolved calculations of laminar helium deflagrations."
 The first definition we consider is the distance between where the temperature is above the upstream temperature aud where the nuclear euerey generatiou attalns its maximi value., The first definition we consider is the distance between where the temperature is above the upstream temperature and where the nuclear energy generation attains its maximum value.
 This width is called the reactive width aud is denoted Syuetear., This width is called the reactive width and is denoted $\delta_{{\rm nuclear}}$.
 The secoud width definition we cousider is the distauce between where the temperature is above the upstream temperature and where the temperature reaches of its downstream value., The second width definition we consider is the distance between where the temperature is above the upstream temperature and where the temperature reaches of its downstream value.
" This width is called the thermal width aud is denoted ?jj,4t"" Lhe", This width is called the thermal width and is denoted $\delta_{{\rm thermal}}$.
 third definition of a deflagratiou's width that we consider is the clistauce between where the composition has its upstream values aud where the clowustreaim composition first reaches its final state., The third definition of a deflagration's width that we consider is the distance between where the composition has its upstream values and where the downstream composition first reaches its final state.
 This width is called the composition width and is denoted Ócomposition:, This width is called the composition width and is denoted $\delta_{{\rm composition}}$.
 For laminar helium deflagrations. the reactive widths jq; are the snallest. widths.," For laminar helium deflagrations, the reactive widths $\delta_{{\rm nuclear}}$ are the smallest widths."
 The thermal widths oiiar are usually slightly larger than the reactive widths dyueear. depeuding ou the upstream thermocdsuzimnic couditious.," The thermal widths $\delta_{{\rm thermal}}$ are usually slightly larger than the reactive widths $\delta_{{\rm
nuclear}}$, depending on the upstream thermodynamic conditions."
 The compositions widtlis 9c;npostion are usually the largest widths. rangiug from being 1.2710 times larger than the thermal widths. depeucling maiuls ou the upstream density.," The compositions widths $\delta_{{\rm composition}}$ are usually the largest widths, ranging from being 1.2–10 times larger than the thermal widths, depending mainly on the upstream density."
 We will use the 9=ὁμιοιοιε 1n our aualysis. but qualitatively similar results are obtained if the dillerences in the widths are taken into account.," We will use the $\delta = \delta_{{\rm nuclear}}$ in our analysis, but qualitatively similar results are obtained if the differences in the widths are taken into account."
 The pressure scale height ofa hydrostatically stratified helium Layer in some cases is comparable to the width ofa helium combustion οί (e.g.. Bildsten 1995). aud should be included in auy length scale comparisons.," The pressure scale height of a hydrostatically stratified helium layer in some cases is comparable to the width of a helium combustion front (e.g., Bildsten 1995), and should be included in any length scale comparisons."
 The pressure scale lieighit is defined as where P? isn the scalar pressure. p isn the mass density.n and g=GÀ//I*H> isn the accelerationn due to gravity.," The pressure scale height is defined as where $P$ is the scalar pressure, $\rho$ is the mass density, and $g =
G M/R^2$ is the acceleration due to gravity."
 For the X-ray. burst case we assumed a AJ=1.1AM.. neutron star with a radius of 2=109 ci. while for the thin shell instability we the aforementioned model from [ben (1977 that is characterized by Al~LOAL.. R~Tx105 em.," For the X-ray burst case we assumed a $M = 1.4 M_{\odot}$ neutron star with a radius of $R=10^6$ cm, while for the thin shell instability we the aforementioned model from Iben (1977) that is characterized by $M \sim
1.0 M_{\odot}$, $R\sim7 \times 10^8$ cm."
 The density in both X-ray burst aud thin shell instability. cases is left as a [ree parameter. permitting a classification of the δα burning regimes as a function of density.," The density in both X-ray burst and thin shell instability cases is left as a free parameter, permitting a classification of the helium burning regimes as a function of density."
 The amplitude of convective velocity fluctuatious ou large scales can be estimated by evaluating, The amplitude of convective velocity fluctuations on large scales can be estimated by evaluating
be significantly allected by charge exchange reactions with hyvelrogen (Péquignot1990:Péquignot.199G):: Consideration of this process would require a knowledge of the ionization state of cloud. which lies beyond the scope of this paper.,"be significantly affected by charge exchange reactions with hydrogen \cite{Peq90,Peq96}: Consideration of this process would require a knowledge of the ionization state of cloud, which lies beyond the scope of this paper."
 Therefore. in our analysis we consider only the case of a primarily neutral.," Therefore, in our analysis we consider only the case of a primarily neutral."
. In fig., In fig.
 6 we plot the population ratios of the erouncl O° fine-structure levels under. various physical conditions., \ref{figure:popOI} we plot the population ratios of the ground $^0$ fine-structure levels under various physical conditions.
 We consider collisions by hydrogen. anc helium atoms. assuming a helium abundance relative to hyerogen of 10 percent (by. number).," We consider collisions by hydrogen and helium atoms, assuming a helium abundance relative to hydrogen of 10 percent (by number)."
 Collisions by helium atoms increases ηςΟρο) by only 5 percent and o(Pu)/o(P5) by 10 percent (reducing to zero close to LYE in the high density limit)., Collisions by helium atoms increases $n({^3\mathrm{P}^e_1})/n({^3\mathrm{P}^e_2})$ by only 5 percent and $n({^3\mathrm{P}^e_0})/n({^3\mathrm{P}^e_2})$ by 10 percent (reducing to zero close to LTE in the high density limit).
 Phe curves for η} corresponding to the inclusion. of the ellects of the CAIBR at z=5 and the UV field of the Galaxy are coincident because the relevantexcitation rates are of the same order Asp-21- , The curves for $n({^3\mathrm{P}^e_1})/n({^3\mathrm{P}^e_2})$ corresponding to the inclusion of the effects of the CMBR at $z=5$ and the UV field of the Galaxy are coincident because the relevantexcitation rates are of the same order $K^{z=5}_{21} \cong \Gamma^{\rmn{G}}_{21}$.
Our results are not. directly. comparable to the work of Péqquignot (1990;1996).. who made assumptions on the ionization state of the gas.," Our results are not directly comparable to the work of Péqquignot \shortcite{Peq90,Peq96}, who made assumptions on the ionization state of the gas."
 We point out. however. the importance of updating the electron. excitationrates emploved in their work - taken from Berrington (1988). - to the more recent caleulations of Bell et ab.," We point out, however, the importance of updating the electron excitationrates employed in their work - taken from Berrington \shortcite{Berrington88} - to the more recent calculations of Bell et al.,"
 since the later's results are substantially lower., since the later's results are substantially lower.
 rpPhe ground state of. the ionH consistsH of. the 3s73p2. 2τιD doublet levels.," The ground state of the $^+$ ion consists of the $^2$ 3p $^2\mathrm{P}^o_{\frac{1}{2},\frac{3}{2}}$ doublet levels."
" Phe energy of the Dine-strueture excited level relatively to the ground. state is 287.24 em|. and. the transition probability is laa—2.1710!s !. Our model ion includesο the three lowest LS terms: 3873pas 2pe6CPU. 38uo 342 dp.OP"" and 3sqo uu: 2]y 7D'. makingF a total of Y levels when the fine-structure is accounted. for."," The energy of the fine-structure excited level relatively to the ground state is 287.24 $^{-1}$, and the transition probability is $A_{\frac{3}{2}\frac{1}{2}}=2.17\ 10^{-4}\ \mathrm{s}^{-1}$ Our model ion includes the three lowest LS terms: $^2$ 3p $^2\mathrm{P}^o$, 3s $^2$ $^4\mathrm{P}^e$ and 3s $^2$ $^2\mathrm{D}^e$ , making a total of 7 levels when the fine-structure is accounted for."
 The energies were taken from. Martin and. Zalubas (1983).., The energies were taken from Martin and Zalubas \shortcite{E_Si_II}. .
" The transition probabilities for the 7D:-OTp'p forbidden transition was taken from. Nusshaumer (1977).. those [or the tp"":CP"" intercombination transitions from Calamai. Smith Bereeson (1993) and those for the 2]pyο allowed: transitions from Nahar (1905)..."," The transition probabilities for the $^2\mathrm{P}^o_{\frac{3}{2}}\rightarrow{^2\mathrm{P}^o_{\frac{1}{2}}}$ forbidden transition was taken from Nussbaumer \shortcite{Nuss77}, , those for the $^4\mathrm{P}^e\rightarrow{^2\mathrm{P}^o}$ intercombination transitions from Calamai, Smith Bergeson \shortcite{CSB1993}
 and those for the $^2\mathrm{D}^e\rightarrow{^2\mathrm{P}^o}$ allowed transitions from Nahar \shortcite{Nahar98}."
 Since the fine-structure levels of Si are too separated apart [rom eachother. the CAIBR will not be an important excitation mechanism.," Since the fine-structure levels of $^+$ are too separated apart from eachother, the CMBR will not be an important excitation mechanism."
 [ven for extremely high redshifts 2o— 5. the excitation rate was lound to be just Aue= 4Collisional processes considered. arecollisions by electrons. (Dufton.&Wineston 1991).. protons(Delv and hydrogen atoms (loucll! 1990)..," Even for extremely high redshifts $z=5$ , the excitation rate was found to be just $K_{\frac{1}{2}\frac{3}{2}}=4.7\ 10^{-15}\ \mathrm{s}^{-1}$ .Collisional processes considered arecollisions by electrons \cite{electron_Si_II}, , protons\cite{p_Si_II} and hydrogen atoms \cite{H0_Si_II}. ."
 In fig., In fig.
using a+ variable parametrization for wea) which is able to accurately describe the expansion history over the full range of redshifts modelled bv the simulations (?)..,using a 4 variable parametrization for $w(a)$ which is able to accurately describe the expansion history over the full range of redshifts modelled by the simulations \citep{Corasaniti:2002vg}.
 The presence of small but appreciable amounts of dark energv at carly times also modcilies the growth. rate of Iluetuations from that expected. in a matter dominated universe and hence changes the shape of the linear theory P(A) from the ACDAL prediction. (2).., The presence of small but appreciable amounts of dark energy at early times also modifies the growth rate of fluctuations from that expected in a matter dominated universe and hence changes the shape of the linear theory $P(k)$ from the $\Lambda$ CDM prediction \citep{2010MNRAS.401.2181J}.
" Phe CNR quintessence mioclel used in this paper has non-negligible amounts of dark οποιον at high redshifts and so could. be classed as an ""early dark model (2)...", The CNR quintessence model used in this paper has non-negligible amounts of dark energy at high redshifts and so could be classed as an early dark model \citep{Doran:2006kp}.
 Xs a result. the linear theory power spectrum is appreciably dilferent from that in à . CDM cosmology. with a broader turnover. (see?.forfutherdetails).," As a result, the linear theory power spectrum is appreciably different from that in a $\Lambda$ CDM cosmology, with a broader turnover, \citep[see][for futher details]{2010MNRAS.401.2181J}."
 Quintessence dark energy. models will not. necessarily agree with observational cata if we adopt the same cosmological parameters as used in the best fitting ACDAL cosmology., Quintessence dark energy models will not necessarily agree with observational data if we adopt the same cosmological parameters as used in the best fitting $\Lambda$ CDM cosmology.
 These best fit parameters were found. using the observational constraints on distances such as the angular diameter distance to last scattering and the sound horizon at this epoch. from the cosmic microwave background. as well as distance measurements from the baryonic acoustic oscillations and “Pvpe Ia supernovae (?)..," These best fit parameters were found using the observational constraints on distances such as the angular diameter distance to last scattering and the sound horizon at this epoch, from the cosmic microwave background, as well as distance measurements from the baryonic acoustic oscillations and Type Ia supernovae \citep{2010MNRAS.401.2181J}."
 In this paper the best fitting cosmological parameters for each. quintessence moclel are used in the N-bocky simulations. as listed in Table 1.," In this paper the best fitting cosmological parameters for each quintessence model are used in the N-body simulations, as listed in Table 1."
 In the left panel of Fig. L..," In the left panel of Fig. \ref{gf},"
 we plot the exact solution for the linear theory growth factor. divided by the scale factor. as a function of redshift together with the fitting formula in Eq. 1..," we plot the exact solution for the linear theory growth factor, divided by the scale factor, as a function of redshift together with the fitting formula in Eq. \ref{linder}."
 Phe 2EXP quintessence model is not plotted in Fig., The 2EXP quintessence model is not plotted in Fig.
 l as the lineargrowth factor for this mocel differs from. ACDAL only at high. redshifts. z> 10.," \ref{gf} as the lineargrowth factor for this model differs from $\Lambda$ CDM only at high redshifts, $z>10$ ."
 7. found that the formula in Eq., \citet{Linder:2005in} found that the formula in Eq.
 1 reproduces the growth factor to better than for ACDAL cosmologics and to ~0.25% for different dynamical quintessence models to the ones considered. in this paper., \ref{linder} reproduces the growth factor to better than for $\Lambda$ CDM cosmologies and to $\sim 0.25$ for different dynamical quintessence models to the ones considered in this paper.
 We have verified that this fitting formula for D is accurate to ~1% for the SUGIUX and 2EXP dark energy models used. in this paper. over a range of redshifts.," We have verified that this fitting formula for $D$ is accurate to $\sim 1\%$ for the SUGRA and 2EXP dark energy models used in this paper, over a range of redshifts."
 Note. in cosmological models which feature non negligible amounts ofdark energy at high redshifts. a further correction factor is needed to this parametrisation (?)..," Note, in cosmological models which feature non negligible amounts ofdark energy at high redshifts, a further correction factor is needed to this parametrisation \citep{Linder:2009kq}."
" Using the parametrization for we) provided by 2? for ""early darkenergv.. 2 proposed a single correction factor which was independent. of redshift."," Using the parametrization for $w(a)$ provided by \citet{Doran:2006kp} for early dark, \citet{Linder:2009kq} proposed a single correction factor which was independent of redshift."
 The CNR model has a high fractional dark energy density at carly timesand as a result we do not expect the linear theory growth to be accurately reproduced by Eq.1. Xs can be seen in Fig., The CNR model has a high fractional dark energy density at early timesand as a result we do not expect the linear theory growth to be accurately reproduced by Eq.\ref{linder}. As can be seen in Fig.
 1. for the CNR model. any correction factor between the fitting formula sugeested by 2? and the exact solution for De would depend on redshift’ and. is not simply a constant.," \ref{gf} for the CNR model, any correction factor between the fitting formula suggested by \citet{Linder:2005in} and the exact solution for $D/a$ would depend on redshift and is not simply a constant."
" In this case. the ""early dark pparametrisation of 7. is not accurate enough to fully describe the dynamics of the CNI quintessence mocel."," In this case, the early dark parametrisation of \citet{Doran:2006kp} is not accurate enough to fully describe the dynamics of the CNR quintessence model."
 This dillerence is 7554 at z—8 for the CNIUmocel. as can be seen in the ratio plot in the leftpanel of Fig 1..," This difference is $\sim$ at $z=8$ for the CNR model, as can be seen in the ratio plot in the leftpanel of Fig \ref{gf}."
 The exact solution for the linear growth rate. f. and the fitting formula in Eq. 1.. ," The exact solution for the linear growth rate, $f$ , and the fitting formula in Eq. \ref{linder}, ,"
f=Ogcies (a). is plotted in," $f = \Omega^{\gamma}_{\rm{m}}(a)$ , is plotted in"
accretion disk theory ancl simulation. in particular on (he magnitude of the Shakura-Sunvaev a parameter.,"accretion disk theory and simulation, in particular on the magnitude of the Shakura-Sunyaev $\alpha$ parameter."
 The reader interested only in the new results of this paper aud not in the backeround fluid mechanical information ean focus on sections ??.. 72? to ??.. and 4..," The reader interested only in the new results of this paper and not in the background fluid mechanical information can focus on sections \ref{ssturb}, , \ref{turbord} to \ref{cor}, and \ref{dis}."
 IIvelrocdvnamic accretion disk mean flows are wiclely believed to be subcritical. i.e.. the viscously relaxed. laminar [low is linearly stable at all Revnolds numbers. at least. locally.," Hydrodynamic accretion disk mean flows are widely believed to be subcritical, i.e., the viscously relaxed laminar flow is linearly stable at all Reynolds numbers, at least locally."
 Furthermore. all experiments and numerical simulations of interest here pertain (o linearly stable flows. ancl we are mostly interested in the local generation of turbulence.," Furthermore, all experiments and numerical simulations of interest here pertain to linearly stable flows, and we are mostly interested in the local generation of turbulence."
 Therefore. I will focus on subcritical flows in this paper.," Therefore, I will focus on subcritical flows in this paper."
 The transition to turbulence is usually rather different in subcritical and supercritical flows., The transition to turbulence is usually rather different in subcritical and supercritical flows.
 Supercritical flows undergo a cascade of precisely defined. bifurcations in parameter space. eventually leading to fully developed turbulence: (hese Gransitions are well documented and reproduced numerically. e.g. for Couette-Tavlor flows (Anderecke/αἱ.1986 and references therein; Marcus 1984a.b)).," Supercritical flows undergo a cascade of precisely defined bifurcations in parameter space, eventually leading to fully developed turbulence; these transitions are well documented and reproduced numerically, e.g. for Couette-Taylor flows \citealt{And86} and references therein; \citealt{Marc84a, Marc84b}) )."
 Turbulence in subcriGical flows. on the contrary. may. abruptly be triggered. most probably by finite amplitude instabilities (DauchotandDaviaud1994): also. the flow apparently evolves from highly intermittent to filly turbulent over a range of Revnolds numbers.," Turbulence in subcritical flows, on the contrary, may abruptly be triggered, most probably by finite amplitude instabilities \citep{Dauch94}; also, the flow apparently evolves from highly intermittent to fully turbulent over a range of Reynolds numbers."
 Furthermore. shear flows can be either (wall-)bounded or free.," Furthermore, shear flows can be either (wall-)bounded or free."
 The distinction refers to the limitation of the flow in the direction where the shear is applied (the transverse or shearwise direction)., The distinction refers to the limitation of the flow in the direction where the shear is applied (the transverse or shearwise direction).
 This difference in boundary conditions influences some of their turbulent properties: indeed. [ree Hows are characterized by a single length-scale. the extent of the shear laver. whereas the distance to the wall introduces a second length scale in wall-bounded flows.," This difference in boundary conditions influences some of their turbulent properties; indeed, free flows are characterized by a single length-scale, the extent of the shear layer, whereas the distance to the wall introduces a second length scale in wall-bounded flows."
 The influence of the other (streamwise ancl spanwise) boundaries is minimized inasmuch as (heir spacing exceeds (he coherence length of the largest turbulent edcdies. ancl as elobally induced perturbations (such as Ekman cireulation) are minimized by appropriate designs of the experimental setups.," The influence of the other (streamwise and spanwise) boundaries is minimized inasmuch as their spacing exceeds the coherence length of the largest turbulent eddies, and as globally induced perturbations (such as Ekman circulation) are minimized by appropriate designs of the experimental setups."
 Shear flows have been actively studied in the past decades. ancl their turbulent properties are now characterized [ον a large variety of settings.," Shear flows have been actively studied in the past decades, and their turbulent properties are now characterized for a large variety of settings."
 In (this section. I will briellv present the subcritical flows which have direct bearing to (he question of hydrodynamic turbulence in accretion disks. namely. plane Couette and [ree shear flows. either rotating or not. flows. aud Ravleigh-stable tidally driven shear flows in (he shearing sheet approximation.," In this section, I will briefly present the subcritical flows which have direct bearing to the question of hydrodynamic turbulence in accretion disks, namely, plane Couette and free shear flows, either rotating or not, Couette-Taylor flows, and Rayleigh-stable tidally driven shear flows in the shearing sheet approximation."
 The first (wo have been studied through bothexperiments and numerical simulations., The first two have been studied through bothexperiments and numerical simulations.
 On the, On the
and 0.29($2 in.,and $0.29^{+0.03}_{-0.16}$ in.
. These measurements are consistent al the confidence level with the neutral C edge al 0.284 keV. Alternatively. we may be detecting either different spectral features or a single feature found at different redshilts in different observations.," These measurements are consistent at the confidence level with the neutral C edge at 0.284 keV. Alternatively, we may be detecting either different spectral features or a single feature found at different redshifts in different observations."
 The enerev of the PSPC feature is consistent with the edge at 0.39 keV and could be produced by an ionizecl eas in [ront of the nuclear source., The energy of the PSPC feature is consistent with the edge at 0.39 keV and could be produced by an ionized gas in front of the nuclear source.
 If the feature seen in BeppoSAX and is produced by the same ion. this would require the gas to be substantially redshifted (2=0.8440.09 in -T and 0.4 in BeppoSAN).," If the feature seen in BeppoSAX and is produced by the same ion, this would require the gas to be substantially redshifted $z=0.34\pm 0.09$ in -T and $z=0.8\pm 0.4$ in BeppoSAX)."
 However. once large redshift are allowed. the PSPC may be seeing moving matter loo. so we cannot exclude that the edge is produced by other elements (e.g. highlv redshilted edge orOvul. or mildly ionized Fe. i.e [roi at 0.23 keV to al 0.39 keV).," However, once large redshift are allowed, the PSPC may be seeing moving matter too, so we cannot exclude that the edge is produced by other elements (e.g. highly redshifted edge or, or mildly ionized Fe, i.e from at 0.23 keV to at 0.39 keV)."
 In anv case this model suggests that the three instruments are seeing a dvnamically variable svstem of clouds and that in the more recent observations (he clouds mav be accelerating toward (he nucleus., In any case this model suggests that the three instruments are seeing a dynamically variable system of clouds and that in the more recent observations the clouds may be accelerating toward the nucleus.
 Outflow velocities of 0.1-0.3 e have been reported in several quasars for the resonant absorption line (Chartas.Brandt&GallagherReeves.O'Brien&Ward.," Outflow velocities of 0.1-0.3 $c$ have been reported in several quasars for the resonant absorption line \citep{chart,reeves}."
 2003)... All of these are high Inminosity objects that are probably radiating al or above their Eddington limit (Ning&Pounds2003).. while M81 is radiating at 0.001-0.02. νι.," All of these are high luminosity objects that are probably radiating at or above their Eddington limit \citep{kp}, while M81 is radiating at 0.001-0.02 $L_{Edd}$."
 A third possibility to fit the absorption features is a gaussian absorption line., A third possibility to fit the absorption features is a gaussian absorption line.
 This model vields by far the best fit to the data., This model yields by far the best fit to the data.
 The line energy is consistent either with the Ίνα line al 0.308 keV or with the Ίνα line at 0.367 in all three instruments. but again the lack of any feature that could be associated with O makes this interpretation doubtful.," The line energy is consistent either with the $\alpha$ line at 0.308 keV or with the $\alpha$ line at 0.367 in all three instruments, but again the lack of any feature that could be associated with O makes this interpretation doubtful."
 The line αἱ 0.57 keV would be consistent with the best fit energy in the PSPC. but it would require a sienificant redshift to explain the leature in (he other instrumentis (z~0.9 in BeppoSAX and z~0.6 in )). implying a very dviamie evolution.," The line at 0.57 keV would be consistent with the best fit energy in the PSPC, but it would require a significant redshift to explain the feature in the other instruments $\sim 0.9$ in BeppoSAX and $\sim0.6$ in ), implying a very dynamic evolution."
 In either case. however. the observed equivalent width (20 eV in the PSPC. 255 eV in SAN/LECS and 318 eV in /ACIS) would require equivalent hydrogen absorbing columns larger than 107? 7? (Nicastro.Fiore&Matt1999) and deep absorption edges for the corresponding ions αἱ 0.74 keV. at 0.39 keV and at 0.49 keV).," In either case, however, the observed equivalent width (20 eV in the PSPC, 255 eV in SAX/LECS and 318 eV in /ACIS) would require equivalent hydrogen absorbing columns larger than $10^{22}$ $^{-2}$ \citep{nica} and deep absorption edges for the corresponding ions at 0.74 keV, at 0.39 keV and at 0.49 keV)."
 The lack of these features seenis (o rule out this interpretation., The lack of these features seems to rule out this interpretation.
 The main problem with the interpretation of the feature as a C edge (either neutral or lonized) is (hat there is no evidence of the corresponding O Ix edge or any other of the edges. expected between 0.5 and 0.8 keV. which we would expect to see in a thermal gas with solar composition.," The main problem with the interpretation of the feature as a C edge (either neutral or ionized) is that there is no evidence of the corresponding O K edge or any other of the edges, expected between 0.5 and 0.8 keV, which we would expect to see in a thermal gas with solar composition."
 The depth of the edge would then require an over-abundance of C with respect to O by a [actor of LO or more., The depth of the edge would then require an over-abundance of C with respect to O by a factor of 10 or more.
 If the O were depleted onto Large dust grains. the O edge could be substantially suppressed.," If the O were depleted onto large dust grains, the O edge could be substantially suppressed."
 Gaskelletal.(2003) argue that large cust erains are common in AGN.," \citet{gaskell}
 argue that large dust grains are common in AGN."
 If the number ratioof C/O is above 1. the dust will be rich in," If the number ratioof C/O is above 1, the dust will be rich in"
Galaxy clusters and. groups are. important both ας cosmological probes. and as laboratories for studying galaxy evolution.,"Galaxy clusters and groups are important both as cosmological probes, and as laboratories for studying galaxy evolution."
 In particular. their deep gravitational potential means that their baryon content should be nearly representative of the Universe as a whole. and that the diffuse gas is at a temperature that is accessible to observation.," In particular, their deep gravitational potential means that their baryon content should be nearly representative of the Universe as a whole, and that the diffuse gas is at a temperature that is accessible to observation."
 As a result. they represent one of the few places where it is possible to study the stars. cold and hot gas. and dark matter in a single system.," As a result, they represent one of the few places where it is possible to study the stars, cold and hot gas, and dark matter in a single system."
 It is now well known that the mass fraction in stars Is not universal. but in general decreases with increasing cluster mass (e.g.22222).," It is now well known that the mass fraction in stars is not universal, but in general decreases with increasing cluster mass \citep[e.g.][]{Eke-groups2,Eke-groups3,RBGMR,Lin03,G+09}."
 On the other hand. the mass fraction of hot gas appears to with mass (e.g.222).," On the other hand, the mass fraction of hot gas appears to with mass \citep[e.g.][]{Vik+05,Sun08,Pratt09}."
" If these systems are closed boxes. then the sum of stellar and gas mass fractions should equal the universal value fj= O,/O,,. where ©), and O,, are the baryon and matter densities. respectively. relative to the critical density."," If these systems are closed boxes, then the sum of stellar and gas mass fractions should equal the universal value $f_b=\Omega_b/\Omega_m$ , where $\Omega_b$ and $\Omega_m$ are the baryon and matter densities, respectively, relative to the critical density."
 Recently ? claimed. this to be the case. in particular arguing that in the lowest mass systems there is significant stellar mass in the intracluster-light and halo of the central galaxy. which completes the baryon fraction.," Recently \citet{GZZ} claimed this to be the case, in particular arguing that in the lowest mass systems there is significant stellar mass in the intracluster-light and halo of the central galaxy, which completes the baryon fraction."
 The interpretation then is that most of the baryons in these low-mass groups have cooled to form stars: his poses a challenge for normal hierarchical models which predict hat more massive clusters are actually built from these groups (2)., The interpretation then is that most of the baryons in these low-mass groups have cooled to form stars; this poses a challenge for normal hierarchical models which predict that more massive clusters are actually built from these groups \citep{BMBE}.
 An important open question then is whether or not there really exists a significant population of groups with AZ.u4fen.c20 yer cent., An important open question then is whether or not there really exists a significant population of groups with $M_{\rm stars}/M_{\rm gas}>20$ per cent.
 The conclusions of ?. depend partly on an extrapolated mean relation for the gas fraction of clusters as a function of mass. aken from ?..," The conclusions of \citet{GZZ} depend partly on an extrapolated mean relation for the gas fraction of clusters as a function of mass, taken from \citet{Vik+05}."
 However. similar conclusions were reached by ?.. based on a small sample of five nearby Abell clusters observed withXMM.," However, similar conclusions were reached by \citet{LLAC}, based on a small sample of five nearby Abell clusters observed with."
 The other possible explanation for the high ratio of stellar-ο-σας mass in groups is that these systems are deficient in X— emitting gas. and associated metals.," The other possible explanation for the high ratio of stellar-to-gas mass in groups is that these systems are deficient in X--ray emitting gas, and associated metals."
 Recent observationsby and ? indicate this is the case: in particular the latter provides, Recent observationsby \citet{G+09} and \citet{RP09} indicate this is the case; in particular the latter provides
and phase distributions. iu conjunction with realistic MBP structuving. our simulations reveal cfficicut. and ubiquitous. mode couversion of longitudinal oscillations iuto their transverse kink counterparts. at twice the driven frequency.,"and phase distributions, in conjunction with realistic MBP structuring, our simulations reveal efficient, and ubiquitous, mode conversion of longitudinal oscillations into their transverse kink counterparts, at twice the driven frequency."
 The simulations clisplayv wave properties cousistent with our observations. and demonstrate how transverse spicule oscillations. with significant amplitudes. readilv exist |in the solar atimosphere. aud therefore have important consequences for cnerey transportation into the solar corona.," The simulations display wave properties consistent with our observations, and demonstrate how transverse spicule oscillations, with significant amplitudes, readily exist in the solar atmosphere, and therefore have important consequences for energy transportation into the solar corona."
 Wo estinate the enerev flux. E. of the observed chromospheric waves using (DePoutienetal.20075). where pis the mass deusity of the flix tube. e is the observed velocity amplitude aud ey is the corresponding Alfvénn speed. defined as c4=BSly with py the maeguetic permeability.," We estimate the energy flux, E, of the observed chromospheric waves using \citep[][]{DeP07b}, where $\rho$ is the mass density of the flux tube, $v$ is the observed velocity amplitude and $v_{A}$ is the corresponding Alfvénn speed, defined as $v_{A} = B/\sqrt\mu_{0}\rho$, with $\mu_{0}$ the magnetic permeability."
 For a mass deusity of p~--z1.3<10“kei °. and a maenetic field strength. Bw~--12 C derived from) a bright uetwork chromospheric model (Vernazzaetal.1981).. coupled with an observed velocity iuplitude of ez15 |. the energy flux in the chromosphere is Ez3«10 ?.," For a mass density of $\rho\approx1.3\times10^{-8}$ $^{-3}$, and a magnetic field strength $B\approx12$ G, derived from a bright network chromospheric model \citep[][]{Ver81}, coupled with an observed velocity amplitude of $v\approx15$ $^{-1}$, the energy flux in the chromosphere is $E\approx3\times10^{5}$ $^{-2}$."
 It is estimated that approximately of the solar surface is covered by MBPs (SáuchezAbueidaetal.2010).. and if each MBP is lhiuked to a single corresponding chromosphlieric Type spicule. it equates to an average global energy 6 660 Wan7.," It is estimated that approximately of the solar surface is covered by MBPs \citep[][]{San10}, and if each MBP is linked to a single corresponding chromospheric Type spicule, it equates to an average global energy of 660 $^{-2}$."
 Current work suggests waves with a- ¢rerey flux zm O00 ? are sufficicutly vigorous to heat he localized corona and launch the solar wind when heir energy is thermalized (Jessetal.2009)., Current work suggests waves with an energy flux $\approx$ 100 $^{-2}$ are sufficiently vigorous to heat the localized corona and launch the solar wind when their energy is thermalized \citep[][]{Jes09}.
. Therefore. a transmission cocficient of 215% through the thin ransition region would provide sufficieut enerev to heat he eutire corona.," Therefore, a transmission coefficient of $\approx$ through the thin transition region would provide sufficient energy to heat the entire corona."
 Reeious ou the solar surface coutaimine Πο] magnetic structures. such as sunspots. pores and aree ΑΠΟ groups. should possess even higher mass densities and magnetic field strenetls. iu addition to a ereater nunibers of spicules.," Regions on the solar surface containing highly magnetic structures, such as sunspots, pores and large MBP groups, should possess even higher mass densities and magnetic field strengths, in addition to a greater numbers of spicules."
 In this reguue. the energy Hux available to heat the corona will be siguificautlv Neher than the müniuun value required to sustain ocalized heating.," In this regime, the energy flux available to heat the corona will be significantly higher than the minimum value required to sustain localized heating."
 We have utilized images of high spatial aud. temporal resolution. obtained with the Rapid Oscillations in the Sol Atinosplere (ROSA) iustruneut at the Dunu Solar Telescope. to reveal how the generation of transverse cls oscillations in Type spicules is a direct result of anode conversion in the lower solu atmosphere.," We have utilized images of high spatial and temporal resolution, obtained with the Rapid Oscillations in the Solar Atmosphere (ROSA) instrument at the Dunn Solar Telescope, to reveal how the generation of transverse kink oscillations in Type spicules is a direct result of mode conversion in the lower solar atmosphere."
 Through comparison of our observations witl advanced svo-dinieusional maeneto-lyvdrodvuamic (ATID) simulations. we show how longitucinal pressure modes in photospheric magnetic bright poiuts (MDDPs). with oxdodieities in the range 130 0 [10 s. funnel upwards hrough the Suus atinmosphere. before couverting iuto οπής imodes at twice the initial frequency. often with auplitudes exceeding 100 km.," Through comparison of our observations with advanced two-dimensional magneto-hydrodynamic (MHD) simulations, we show how longitudinal pressure modes in photospheric magnetic bright points (MBPs), with periodicities in the range 130 – 440 s, funnel upwards through the Sun's atmosphere, before converting into kink modes at twice the initial frequency, often with amplitudes exceeding 400 km."
 We conclude that the uode conversion and period modification is a direct consequence of the 90 degree phase shift eunconrpassiug opposite sides of the photospheric driver., We conclude that the mode conversion and period modification is a direct consequence of the 90 degree phase shift encompassing opposite sides of the photospheric driver.
 This is the first observational evidence of this ποαπάμα occurring in the solar atmosphere., This is the first observational evidence of this mechanism occurring in the solar atmosphere.
 Euergv flux estimates for hese oscillations indicate that the waves are sufficieutlv cherectic to accelerate the solar wind and heat the quiet corona to its multianillion degree temperatures., Energy flux estimates for these oscillations indicate that the waves are sufficiently energetic to accelerate the solar wind and heat the quiet corona to its multi-million degree temperatures.
" The naming of transverse oscillations observed iu solar structures remains πο] coutroversial. witlidefinitions revolving around “Alfvén”. ~Alfvénic™. and ""unagneto-sonie kink terminology."," The naming of transverse oscillations observed in solar structures remains highly controversial, withdefinitions revolving around “Alfvénn”, ”, and “magneto-sonic kink” terminology."
 While our observations demonstrate signatures consistent witli previous studies ou waves (c.g.DePontieuctal.2007b:MclIutoshet 2011).. we have deliherately chosen to describe the observed periodic motions simply as trausverse kink waves.," While our observations demonstrate signatures consistent with previous studies on waves \citep[e.g.][]{DeP07b, McI11}, we have deliberately chosen to describe the observed periodic motions simply as transverse kink waves."
 This is currently the inost unopposed description of such wave phenomena. aud avoids the potential pitfalls of what is a rapidly developing area within solar plivsics.," This is currently the most unopposed description of such wave phenomena, and avoids the potential pitfalls of what is a rapidly developing area within solar physics."
 DDJ wishes to thank STFC for the award of a Post-Doctoral Fellowship., DBJ wishes to thank STFC for the award of a Post-Doctoral Fellowship.
 DJC thanks CSUN for start-up fuudiueg related to this project., DJC thanks CSUN for start-up funding related to this project.
 PITIS is exateful to NIDEL for a PhD studeutslip., PHK is grateful to NIDEL for a PhD studentship.
 Solar Plivsies research at QUD is supported by STFC., Solar Physics research at QUB is supported by STFC.
 The ROSA project is supported by EOARD. (ROSA)..., The ROSA project is supported by EOARD. .
We acknowledge the inspiring ideas of D. Balick on which this work is based. and the anonymous referee for his/her positive and. helpful comments.,"We acknowledge the inspiring ideas of B. Balick on which this work is based, and the anonymous referee for his/her positive and helpful comments."
 We thank the Instituto cle sica de Canarias for the use of the PC cluster beoiac," We thank the Instituto de sica de Canarias for the use of the PC cluster “beoiac""."
 The work of DRG is supported by the Brazilian Agcnev EAPIESDP. (0X/11837-0)., The work of DRG is supported by the Brazilian Agency FAPESP (04/11837-0).
 BE would like to thank the EXPESP for the visiting erant (03/09692-0)., BE would like to thank the FAPESP for the visiting grant (03/09692-0).
 We also acknowledge the partial support of the Spanish Ministry. of Science ancl Technology (AYA 2002-0883), We also acknowledge the partial support of the Spanish Ministry of Science and Technology (AYA 2002-0883).
2004).,.
". This value is smaller than the value (IL, = 72) obtained bv Freedmanetal.(2001) before removal of the intrinsic components.", This value is smaller than the value $_{o}$ = 72) obtained by \citet{fre01} before removal of the intrinsic components.
 In most respects the DIR model is perfectly. compatible with the standard Bie Bane model of the Universe., In most respects the DIR model is perfectly compatible with the standard Big Bang model of the Universe.
 H differs mainly in the way galaxies are born and the elaim that in this model at least the radio galaxies pass through an initial short-lived AGN period (105 vrs) in which their redshifts contain an intrinsic component (hat quickly disappears., It differs mainly in the way galaxies are born and the claim that in this model at least the radio galaxies pass through an initial short-lived AGN period $^{8}$ yrs) in which their redshifts contain an intrinsic component that quickly disappears.
 After that. as they evolve through the next 1! Vears they can be used as thev are today. to study cosmology.," After that, as they evolve through the next $^{10}$ years they can be used as they are today, to study cosmology."
 Although there is now a considerable amount of evidence supporting the DUR model. (here are also some well-known arguments against this model that have been raised by those who support the CR. model (e.g. the Lyman forest. lensing by intervening galaxies. etc.).," Although there is now a considerable amount of evidence supporting the DIR model, there are also some well-known arguments against this model that have been raised by those who support the CR model (e.g. the Lyman forest, lensing by intervening galaxies, etc.)."
 An explanation of (hese arguments in the DIR model can be found in the Discussion section of a previous paper (Dell2004)., An explanation of these arguments in the DIR model can be found in the Discussion section of a previous paper \citep{bel04}.
". In the CR. model the location of high-redshift AGN galaxies. (quasars) on a logz-am,. plot can be explained by the presence of a non-thermal component superimposed on (heir optical huminosity.", In the CR model the location of high-redshift AGN galaxies (quasars) on a $z$ $_{v}$ plot can be explained by the presence of a non-thermal component superimposed on their optical luminosity.
 In the DIR model their location on this plot is explained bv the presence of a non-cosmological redshilt component superimposed on (heir redshift., In the DIR model their location on this plot is explained by the presence of a non-cosmological redshift component superimposed on their redshift.
" This paper uses an updated logz-m, plot containing over 100.000 AGN &alaxies to compare the most luminous radio galaxies and first-rankecl cluster galaxies at each redshift to the high luminosity edge ol the AGN galaxy. distribution in an attempt to see which model (CR or DIR. model) can best explain the data."," This paper uses an updated $z$ $_{v}$ plot containing over 100,000 AGN galaxies to compare the most luminous radio galaxies and first-ranked cluster galaxies at each redshift to the high luminosity edge of the AGN galaxy distribution in an attempt to see which model (CR or DIR model) can best explain the data."
 In this paper the standard candle (constant. Iuminositv) slope is used as a relerence {ο make Iuminositv comparisons at a given redshilt., In this paper the standard candle (constant luminosity) slope is used as a reference to make luminosity comparisons at a given redshift.
 This is shown as a dashed line in Fige 1 and a solid line in Fige 2., This is shown as a dashed line in Fig 1 and a solid line in Fig 2.
 Luminosity increases to the left., Luminosity increases to the left.
" A logz-m, plot for those radio sources with measured redshifts that were detected in the 1 Jv radio survey. (Stickeletal.1994). is presented in Fig 1.", A $_{v}$ plot for those radio sources with measured redshifts that were detected in the 1 Jy radio survey \citep{sti94} is presented in Fig 1.
 The quasars are plotted as filled circles and the radio galaxies as open squares., The quasars are plotted as filled circles and the radio galaxies as open squares.
 As discussed above. in the DIR. model the radio galaxies ave the objects that high-redshift quasars ancl other AGN galaxies evolve into when their intrinsic redshift component has largely disappeared.," As discussed above, in the DIR model the radio galaxies are the objects that high-redshift quasars and other AGN galaxies evolve into when their intrinsic redshift component has largely disappeared."
 In Fig 1. lirst-rankecl cluster galaxies (Sandage1972a:Kristianetal.1978). are indicated bv the dashed line.," In Fig 1, first-ranked cluster galaxies \citep{san72a,kri78} are indicated by the dashed line."
 The most luminous radio galaxies. like first-rankecl cluster galaxies. are clearly good. standard candles to large cosmological distances. and their redshifts must then be cosmological. as expected in both the CR and DUR models since anv intrinsic redshift component. will have almost completely disappeared.," The most luminous radio galaxies, like first-ranked cluster galaxies, are clearly good standard candles to large cosmological distances, and their redshifts must then be cosmological, as expected in both the CR and DIR models since any intrinsic redshift component will have almost completely disappeared."
 All the sources listed as quasars and active galaxies in the updated: Véron-Cettv/Véron, All the sources listed as quasars and active galaxies in the updated $\acute{e}$ $\acute{e}$ ron
 , 
"increases 1, by a factor of 10.",increases $\nut$ by a factor of 10.
 This is consistent with the classical theory of a boundary layer. according to which the thickness of the boundary layer is inversely proportional to the square root of the Reynolds number. VRe~|jvh (see.forexample. 1962).," This is consistent with the classical theory of a boundary layer, according to which the thickness of the boundary layer is inversely proportional to the square root of the Reynolds number, $\sqrt{Re} \sim 1/\sqrt \nut$ \citep[see, for example, ][]{sip62}."
. The increase of viscosity also leads to a growth of the gas temperature. resulting from a balance between turbulent viscous friction and radiative cooling.," The increase of viscosity also leads to a growth of the gas temperature, resulting from a balance between turbulent viscous friction and radiative cooling."
 One can see that the temperature achieves its maximal value in the middle of the boundary layer. where the velocity gradient and therefore the heating rate are maximum.," One can see that the temperature achieves its maximal value in the middle of the boundary layer, where the velocity gradient and therefore the heating rate are maximum."
 Note. that the solution for our test case looks similar to the spreading laver model for the white dwarf case.," Note, that the solution for our test case looks similar to the spreading layer model for the white dwarf case."
 We find that the Z-component of the gas velocity Vz=10° em ∣ is very close to that obtained by Piro&Bildsten(2004)., We find that the $Z$ -component of the gas velocity $V_Z=10^6$ cm $^{-1}$ is very close to that obtained by \citet{pb04}.
. Unfortunately. we cannot compare other quantities because the values in the spreading layer model are averaged along the A-direction.," Unfortunately, we cannot compare other quantities because the values in the spreading layer model are averaged along the $R$ -direction."
" Let us now consider a realistic neutron star mass. Ma,=I.1M.."," Let us now consider a realistic neutron star mass, $M_{\rm
star}=1.4\msun$."
 We find that. in order to balance the gravity force by the gas pressure gradient near the stellar surface. the gas temperature must attain a value of 3x10'7 ΚΚ. which is unrealistic.," We find that, in order to balance the gravity force by the gas pressure gradient near the stellar surface, the gas temperature must attain a value of $3\times 10^{12}$ K, which is unrealistic."
 Therefore. in the case of a non-rotating star. the only force which can work against gravity is the radiation pressure gradient.," Therefore, in the case of a non-rotating star, the only force which can work against gravity is the radiation pressure gradient."
" Indeed. if one takes the gas temperature to be 7=3xLOS K and =0.3 g em. the radiation pressure force becomes comparable to the gravitational force ΜΙΑ<F1,Cup(ορ). where cg is the Boltzmann constant)."," Indeed, if one takes the gas temperature to be $T= 3 \times 10^8$ K and $\rho \gtrsim 0.3$ g $^{-3}$, the radiation pressure force becomes comparable to the gravitational force $GM_{\rm star}/R^2 \lesssim \sigma_{\rm SB} T_{\rm
star}^4/(c \rho)$, where $\sigma_{\rm SB}$ is the Stefan-Boltzmann constant)."
" Using standard disk theory. we estimate the turbulent viscosity near the stellar surface v,=eHc,LOM ene !; where aqua,=0.01 is a viscosity parameter. H20.01Ruy. and c,210 em |,"," Using standard disk theory, we estimate the turbulent viscosity near the stellar surface $\nut =\alpha_{\rm d} H c_{\rm s}\simeq
10^{10}$ $^{2}$ $^{-1}$, where $\alpha_{\rm disk} \simeq 0.01$ is a viscosity parameter, $H\simeq 0.01 R_{\rm star}$, and $c_{\rm s}\simeq
10^8$ cm $^{-1}$."
" Note. that the radiative viscosity 1,=docuPmκρ)10% ene ! (where nn, is à proton mass) is much smaller than the turbulent one. and can be neglected."," Note, that the radiative viscosity $\nu_{\rm r}=4
\sigma_{\rm SB} T^4 m_{\rm p}/(\kappa c^2 \rho) \simeq 10^8$ $^{2}$ $^{-1}$ (where $m_{\rm p}$ is a proton mass) is much smaller than the turbulent one, and can be neglected."
 We tind that the description of butter zones must be modified in the radiation pressure dominated case., We find that the description of buffer zones must be modified in the radiation pressure dominated case.
 For simplicity we use a similar density profile in the disk buffer zone as it was in the gas pressure dominated case., For simplicity we use a similar density profile in the disk buffer zone as it was in the gas pressure dominated case.
" However. we now takepq(0)=4 zem and fix the gas temperature μις to avoid large radiation pressure gradients and therefore the generation of large velocities. which lead to strongly non-stationary behavior in the butter zone and eventually to numerical instability,"," However, we now take $\rho_{\rm disk}(0)=4 $ g $^{-3}$ and fix the gas temperature $T_{\rm disk}$ to avoid large radiation pressure gradients and therefore the generation of large velocities, which lead to strongly non-stationary behavior in the buffer zone and eventually to numerical instability."
 Also. we use a softer condition for the Z-velocity by assuming vanishing first derivatives.," Also, we use a softer condition for the $Z$ -velocity by assuming vanishing first derivatives."
 Thus. we set In the course of the calculation we tind that inside the domain cold low-density patches surrounded by denser hot gas appear sporadically.," Thus, we set In the course of the calculation we find that inside the domain cold low-density patches surrounded by denser hot gas appear sporadically."
 Such patches are dispersed due to motion of the gas from the core of the patch outward through the radiation pressure gradient., Such patches are dispersed due to motion of the gas from the core of the patch outward through the radiation pressure gradient.
 However. if we were to fix the Z and & velocities in the star buffer zone to zero. such a patch cannot disappear. while the," However, if we were to fix the $Z$ and $R$ velocities in the star buffer zone to zero, such a patch cannot disappear, while the"
discuss in detail the exact imechanisia for expelling mass from the system diving MT. aud do not propose what the exact value of the efficiency should be. but rather study the consequences of nou-conservationt during MT with au adopted conservation factor.,"discuss in detail the exact mechanism for expelling mass from the system during MT, and do not propose what the exact value of the efficiency should be, but rather study the consequences of non-conservation during MT with an adopted conservation factor."
 We may argue however that MT is rather Likely to be nonu-conservative. especially during the TPTAIT phase.," We may argue however that MT is rather likely to be non-conservative, especially during the TTMT phase."
 The strongest arguineut for this is due to the very high AIT rate expected: some of the binary systems. at the peak of TINT. could have MT rates above the Eddington Πιτ (~10CAZor. for MS accretors in the mass range we considered).," The strongest argument for this is due to the very high MT rate expected: some of the binary systems, at the peak of TTMT, could have MT rates above the Eddington limit $\sim 10^{-3}\,M_\odot \mathrm{yr}^{-1}$ for MS accretors in the mass range we considered)."
 Observationally.+ though. it is hard to make any coustraiut: the TTAIT phase is of course very short lived.," Observationally, though, it is hard to make any constraint: the TTMT phase is of course very short lived."
 We note however that among observed biuaries undergoing ALT with giant donors. e.g.. in svinbiotic stars. there are observed systems showing powerful jets clearly. the MT is not filly conservative there (e.@.. in svinhiotic stars like CI Cveui (?).. CIT ονο (233). even though the AIT proceeds there on a nuclear timescale.," We note however that among observed binaries undergoing MT with giant donors, e.g., in symbiotic stars, there are observed systems showing powerful jets — clearly, the MT is not fully conservative there (e.g., in symbiotic stars like CI Cygni \citep{CICYG}, CH Cygni \citep{CHCyg}) ), even though the MT proceeds there on a nuclear timescale."
 Young pre-MS stars are the only observed. non-conipact objects that accrete at rates almost comparable to our T'TMT rates. aud are known to have jets as well e.c.. IIT 30 (25.," Young pre-MS stars are the only observed non-compact objects that accrete at rates almost comparable to our TTMT rates, and are known to have jets as well e.g., HH 30 \citep{HH30}."
 Thus our assmuptiou that bipolar re-ciission from the accretor drives non-conservation in MT. aud therefore the specific augular moimeutun of the lost material is that of the accretor. scemis at least im principle quite reasonable.," Thus our assumption that bipolar re-emission from the accretor drives non-conservation in MT, and therefore the specific angular momentum of the lost material is that of the accretor, seems at least in principle quite reasonable."
 Iu the coutest of earlier studies. it is important to note that the progenitor donors with ALz2M. often considered previously would uot be able to form à WD inuncdiately following au episode of MT they would instead formi a low-mass Πο star with a lifetime of up to O5GCar C8). relevant to the age difference of mma of the solutions described iu table 5 of ?..," In the context of earlier studies, it is important to note that the progenitor donors with $M \gtrsim 2\,M_\odot$ often considered previously would not be able to form a WD immediately following an episode of MT — they would instead form a low-mass He star with a lifetime of up to Gyr \citep{Yungelson08}, relevant to the age difference of many of the solutions described in table 5 of \cite{Sluys2006}."
 Tn this study we find that iu the low-mass reeine the stable AIT | CE channel. with au initially 1.0—1.5AZ... donor. provides a natural explanation for all but one of the observed doubleavhite-dwarf binaries for which the inferred older component is below 0.16AL...," In this study we find that in the low-mass regime the stable MT + CE channel, with an initially $1.0-1.3\,M_\odot$ donor, provides a natural explanation for all but one of the observed double-white-dwarf binaries for which the inferred older component is below $0.46\,M_\odot$."
 Thus. svstems such as H1115|116. in which the vouueer conrpanuion is the more massive. can casily form as a result of the orbital expansion during the initial phase of RLOF. followed by a CE.," Thus, systems such as 1115+116, in which the younger companion is the more massive, can easily form as a result of the orbital expansion during the initial phase of RLOF, followed by a CE."
 The relatively small delay between the two phases of mass loss from the system. inferred from observations. is also easily justified by the accelerated evolution of the accretor as a result of the first AIT episode.," The relatively small delay between the two phases of mass loss from the system, inferred from observations, is also easily justified by the accelerated evolution of the accretor as a result of the first MT episode."
 Of course. a stroug determination of the relative ages in a DWD is a difficult task. further confounded by the possibility of rejuvenation.," Of course, a strong determination of the relative ages in a DWD is a difficult task, further confounded by the possibility of rejuvenation."
 The system 1110112361 presents a clear challenge in that its older compoucut is the more massive. vet it cannot be explained as the product of a double common euvelope without resorting to unphysical efficiencies iu the first phase of mass loss.," The system 1101+364 presents a clear challenge in that its older component is the more massive, yet it cannot be explained as the product of a double common envelope without resorting to unphysical efficiencies in the first phase of mass loss."
 One possibility may be sole iutermediary process. such as that sugeested by 7.. in which the envelope of the initial primary ds lost without significant orbital shrinkage: ?— suggests this stage iav be understood as a phase of very rapid wind loss.," One possibility may be some intermediary process, such as that suggested by \cite{Nelemans00}, in which the envelope of the initial primary is lost without significant orbital shrinkage; \cite{Sluys2006} suggests this stage may be understood as a phase of very rapid wind loss."
 However. au obvious mechanisin for this remains lacking.," However, an obvious mechanism for this remains lacking."
 It is also possible that the formation of a siuall circunibinary disk in the first plase of mass transfer wav allow for some shrinkage of the orbit. however there is a lack of evidence at the present moment to support such a lbypothesis.," It is also possible that the formation of a small circumbinary disk in the first phase of mass transfer may allow for some shrinkage of the orbit, however there is a lack of evidence at the present moment to support such a hypothesis."
 Perhaps the most likely explanation. eiven the small secondary mass in II101|361 aud herefore its thicker lyvdrogen envelope. is that the resulting uncertainty iu the cooling tracks means that he 0.29AL. companion is older than previously thought.," Perhaps the most likely explanation, given the small secondary mass in 1101+364 and therefore its thicker hydrogen envelope, is that the resulting uncertainty in the cooling tracks means that the $0.29\,M_\odot$ companion is older than previously thought."
" and the host galaxy metallicity, mass) are kinematics)well constrained (Chen(luminosity,&Prochaska2000;sonetal.","kinematics) and the host galaxy (luminosity, metallicity, mass) are well constrained \citep{cp00,jenkins05,tripp05,prochaska06,cooksey08,lehner09,tumlinson11}."
" To understand the implications for the statistical 2011)..evolution of LLSs in the context of galaxy evolution it is necessary to better understand the gas-galaxy relationship, including the frequency with which LLSs trace infall onto versus outflows from galaxies."," To understand the implications for the statistical evolution of LLSs in the context of galaxy evolution it is necessary to better understand the gas-galaxy relationship, including the frequency with which LLSs trace infall onto versus outflows from galaxies."
" Here we use observations from the Cosmic Origins Spectrograph (COS), with additional ground-based spectroscopic and imaging observations, to analyze a LLS at z~0.274 along the sight line to the UV-bright QSO PG1630--377 (zem= 1.476)."," Here we use observations from the Cosmic Origins Spectrograph (COS), with additional ground-based spectroscopic and imaging observations, to analyze a LLS at $z\sim0.274$ along the sight line to the UV-bright QSO PG1630+377 $z_{\rm em}=1.476$ )."
" We demonstrate the LLS has low metallicity and use Large Binocular Telescope (LBT) imaging(8 2))and Keck spectroscopy of galaxies in the QSO field to identify a near-solar metallicity, 0.3 Τι galaxy at virtually the same redshift "," We demonstrate the LLS has low metallicity \ref{sec:absorption}) ) and use Large Binocular Telescope (LBT) imaging and Keck spectroscopy of galaxies in the QSO field to identify a near-solar metallicity, 0.3 $_{*}$ galaxy at virtually the same redshift \ref{sec:galenv}) )."
We discuss this gas in the context of cold mode accretion(83)). models in 4 and summarize our results in ?7.," We discuss this gas in the context of cold mode accretion models in \ref{sec:origins}
 and summarize our results in \ref{sec-sum}."
" The UV spectra of PG1630+377 were obtained with COS on-board the (PID 11741, PI Tripp) using the G130M (1150-1450 À)) and G160M (1405-1775 À)) gratings."," The UV spectra of PG1630+377 were obtained with COS on-board the (PID 11741, PI Tripp) using the G130M (1150–1450 ) and G160M (1405–1775 ) gratings."
" The exposure times for these two configurations were 23.0 and 14.3 ks, respectively, giving S/N up to 30-40 per resolution element at unabsorbed wavelengths."," The exposure times for these two configurations were 23.0 and 14.3 ks, respectively, giving S/N up to 30–40 per resolution element at unabsorbed wavelengths."
 The excellent quality of the UV spectra is displayed in Figures 1 and 2.., The excellent quality of the UV spectra is displayed in Figures \ref{fig:big} and \ref{fig:profile}.
 The data were processed using CALCOS and coadded following Meiringetal.(2011)., The data were processed using CALCOS (v2.11b) and coadded following \citet{meiring11}.
". The (v2.11b)lower-right panel of Figure 1 shows a small portion of the spectrum obtained from COS, highlighting the presence of a LLS at z~0.274."," The lower-right panel of Figure \ref{fig:big} shows a small portion of the spectrum obtained from COS, highlighting the presence of a LLS at $z\sim0.274$."
" Optical spectra were obtained with HIRES at the Keck Observatory on 26 March 2010 and cover AA2796,2803 at z~0.274."," Optical spectra were obtained with HIRES at the Keck Observatory on 26 March 2010 and cover $\lambda\lambda2796,2803$ at $z\sim0.274$."
" We acquired 3 HIRES exposures totaling 2.2 ks under good conditions with the blue cross-disperser, the C1 decker wwidth giving FWHM 5»6 s-1)), and the CCD mosaic."," We acquired 3 HIRES exposures totaling 2.2 ks under good conditions with the blue cross-disperser, the C1 decker width giving FWHM $\approx 6$ ), and the CCD mosaic."
" The data were processed using HIRedux in Figure 2 shows the normalized absorption profiles of and metal ions as a function of ((leftvelocity panels)relative to zaps=0.27395,panels) the centroid of aabsorption."," The data were processed using HIRedux in Figure \ref{fig:profile} shows the normalized absorption profiles of ) and metal ions ) as a function of velocity relative to $z_{\rm abs}=\zone$, the centroid of absorption."
" At v—0 wwe detectL,IIL, IIL, weak aabsorption (W,(2796)=59+4 m), and possiblyVIL."," At $v=0$ we detect, , weak absorption $W_r(2796)=59 \pm 4$ ), and possibly."
 We refer to this as the strong due to the strength of the aabsorption., We refer to this as the strong due to the strength of the absorption.
" Near v=+50 ((z= 0.27416) we detect L,IIL,IIL, and VL, but noII."," Near $v=+50$ $z=\ztwo$ ) we detect , and , but no."
. We refer to this as the weak component., We refer to this as the weak component.
" In Figure 3 we show the apparent column density profiles, Να{υ) (Savage&Sembach1991),, of selected species."," In Figure \ref{fig:nav} we show the apparent column density profiles, $N_a(v)$ \citep[][]{savage91}, of selected species."
 'This figure demonstrates that the ffollows the intermediate ion aand iin the weak component well and that all of the observedccould be due to the weak component., This figure demonstrates that the follows the intermediate ion and in the weak component well and that all of the observedcould be due to the weak component.
 'The column densities and limitsfor selected species can, The column densities and limitsfor selected species can
AGB stars. (wo unresolved rreeions coincident with strong Ha emission. and (wo objects that are probably foreground ejants.,"AGB stars, two unresolved regions coincident with strong $\alpha$ emission, and two objects that are probably foreground giants."
 It is likely that many of the objects not detected in the optical are dust enshrouded AGB stars. since only the brightest and bluest of these objects are detected optically and the majority of (hese are optically classified as AGBs (see the discussion in 8??)).," It is likely that many of the objects not detected in the optical are dust enshrouded AGB stars, since only the brightest and bluest of these objects are detected optically and the majority of these are optically classified as AGBs (see the discussion in \ref{mass_loss}) )."
 Due to the micro-Jyv point source sensitivity of IRAC with our integration time. we should be able to detect the entire population of the AGB stars in WLM. excluding only (hose objects Chat were not detected due to crowding.," Due to the micro-Jy point source sensitivity of IRAC with our integration time, we should be able to detect the entire population of the AGB stars in WLM, excluding only those objects that were not detected due to crowding."
 However. uniquely determining the true stellar tvpes of the detected objects can be challenging. as (he spectral energy. distributions (SEDs) of these stellar (vpes. with the exception of unresolved rregions. are indistinguishable [rom one another.," However, uniquely determining the true stellar types of the detected objects can be challenging, as the spectral energy distributions (SEDs) of these stellar types, with the exception of unresolved regions, are indistinguishable from one another."
 Figure 9 shows the IRAC SEDs for the all of the AGB stars. RSGs. and blue objects detected in all four IRAC bands ancl the LGGS V and I images.," Figure \ref{SED} shows the IRAC SEDs for the all of the AGB stars, RSGs, and blue objects detected in all four IRAC bands and the LGGS V and I images."
 This plot shows the similarity between the infrared f[Iuxes of AGDs and RSCs., This plot shows the similarity between the infrared fluxes of AGBs and RSGs.
 Although it is challenging to distinguish between AGB stars and RSGs using IRAC data alone. the combination of optical and data provide a powerful method of detecting candidate AGB stars in galaxies past the immediate vicinity of the Milkv. Was. since RSGs are generally easily ilentilied in optical CMDSs.," Although it is challenging to distinguish between AGB stars and RSGs using IRAC data alone, the combination of optical and near-IR data provide a powerful method of detecting candidate AGB stars in galaxies past the immediate vicinity of the Milky Way, since RSGs are generally easily identified in optical CMDs."
 One of the goals of our study is to determine the fraction of AGB stars seen in the IR that were not detected in the optical due to extinction by cireumstellar material., One of the goals of our study is to determine the fraction of AGB stars seen in the IR that were not detected in the optical due to extinction by circumstellar material.
 To estimate (his fraction. we make the approximation (hat all objects above the TRGB in the IB. CMD ave AGB stars.," To estimate this fraction, we make the approximation that all objects above the TRGB in the IR CMD are AGB stars."
 While there is contamination above the TRGB from RSCs and blue objects. the number of these objects detected in the optical contribute <LO%. of the population above ihe TRGB in the IR.," While there is contamination above the TRGB from RSGs and blue objects, the number of these objects detected in the optical contribute $\lesssim$ of the population above the TRGB in the IR."
 Of the 691 stars brighter than the TRGB in the IR. are not detected in V and I. There are also 29 objects brighter than the TRGB in the IR CMD. but le within region (d)," Of the 691 stars brighter than the TRGB in the IR, are not detected in V and I. There are also 29 objects brighter than the TRGB in the IR CMD, but lie within region (d)"
The character of the SAAO instrumental svstem plavs a Κον role in our work (see 83).,The character of the SAAO instrumental system plays a key role in our work (see 3).
 That svstem has a history of vielding stable translormations to the Cousins svstem (see. lor example. 32 of Joneretal. 2006)).," That system has a history of yielding stable transformations to the Cousins system (see, for example, 2 of \citealt{jtlw06}) )."
 Moreover. the SAAO svstem has been used extensively to establish Cousins standards in the southern hemisphere (see. [or example. MenziesNilkennyetal.1905:Koen 2002)).," Moreover, the SAAO system has been used extensively to establish Cousins standards in the southern hemisphere (see, for example, \citealt{mc89,kw98,kk02}) )."
 For these reasons (see also 83.2). the SAAO system is regarded here as (he most authoritative svstem of its kind.," For these reasons (see also 3.2), the SAAO system is regarded here as the most authoritative system of its kind."
 All but one of the papers produced by the projects are listed in Table I., All but one of the papers produced by the projects are listed in Table 1.
 A paper by Tavlor&Joner(1996). that compares the Cousins ancl Landolt(1983) systems is nol listed in the table because it contains no cluster photometry., A paper by \citet{tj96} that compares the Cousins and \citet{l83} systems is not listed in the table because it contains no cluster photometry.
 However. that paper should be regarded as part of the project (see 83.1).," However, that paper should be regarded as part of the project (see 3.1)."
 We also note that the first SAAO paper listed in Table 1 has been cited several times (see. lor example. Twarogοἱal.2009. and 2011)).," We also note that the first SAAO paper listed in Table 1 has been cited several times (see, for example, \citealt{tate09} and \citealt{clhw11}) )."
 However. TJJ has not been cited often. and it appears (hat its existence is nol widely known.," However, TJJ has not been cited often, and it appears that its existence is not widely known."
 Potential users of project data are therefore invited to adopt TJJ data for the Ivacles. M67. Coma. and/or NGC 752.," Potential users of project data are therefore invited to adopt TJJ data for the Hyades, M67, Coma, and/or NGC 752."
 IE HEvades data are used. the data specified in Table 2 may be added.," If Hyades data are used, the data specified in Table 2 may be added."
 Before that is done. however. users should make a point of reading ihe Table 2 footnotes.," Before that is done, however, users should make a point of reading the Table 2 footnotes."
 A note is needed about project papers containing Ivaces data., A note is needed about project papers containing Hyades data.
 Users who do not require Ivades values of V—R can limit their attention to Table 2 and CDS data files from Those who wish to include ILvades values of V.—R should consult Tavlor (1985).. Joneretal. (2006).. and Joneretal.(2008) and the footnotes given in Table 1 to the entries [ον (hose papers.," Users who do not require Hyades values of $V-R$ can limit their attention to Table 2 and CDS data files from Those who wish to include Hyades values of $V-R$ should consult \citet{tj85}, \citet{jtlw06}, and \citet{jtlw08} and the footnotes given in Table 1 to the entries for those papers."
 One aim of the projects has been to determine exact relationships between their published data and the Cousins. Johnson. and Strómnmegren-2 svstems.," One aim of the projects has been to determine exact relationships between their published data and the Cousins, Johnson, and $\beta$ systems."
 A second aim (as noted in 81) has been to establish zero points that are known to be uniform from cluster to cluster., A second aim (as noted in 1) has been to establish zero points that are known to be uniform from cluster to cluster.
 To support the latter aim. a number of nights have been used to compare (vo or three clusters directly.," To support the latter aim, a number of nights have been used to compare two or three clusters directly."
" These ""Sturch comparisons” imitate a procedure applied by Sturch(1972.1973)."," These “Sturch comparisons” imitate a procedure applied by \citet{s72,s73}."
". In addition. zero point. accuracy has been assessed [rom statistical comparisons of independent data sets for ""SCIDS [or short)."," In addition, zero point accuracy has been assessed from statistical comparisons of independent data sets (or “SCIDS” for short)."
" The mean residuals obtained [rom SCIDS satislv the “FATT (7few millimag"") standardthat is. thev are typically a few manag in size."," The mean residuals obtained from SCIDS satisfy the so-called “FM” (“few millimag”) standard–that is, they are typically a few mmag in size."
 As explained in, As explained in
in the subroutine available at the website described given in Section 6..,in the subroutine available at the website described given in Section \ref{sec:conc}.
 As au example. the coellicieuts for filter FOOGW are reported in Table 2. (coelficieuts for all the teu filters are shown in Table A in the electronic version of the paper).," As an example, the coefficients for filter F606W are reported in Table \ref{tab2} (coefficients for all the ten filters are shown in Table A in the electronic version of the paper)."
 The next step iu the procedure was to examine the residuals from the polynomial solution., The next step in the procedure was to examine the residuals from the polynomial solution.
 To do this we trausforiued the imaster-fraiue position for each star back iuto the raw frame ol each exposure using a conformal linear land the inverse distortion solution., To do this we transformed the master-frame position for each star back into the raw frame of each exposure using a conformal linear and the inverse distortion solution.
 The residual was then the difference between where the star was observed in the [rame Creag.ga) aud where it should have been in that [ranue. according to the master (rame.," The residual was then the difference between where the star was observed in the frame $(x_{\rm raw},y_{\rm raw})$ and where it should have been in that frame, according to the master frame."
 Iu Figure 1 we plot this residual as a function of raw chip coordinate., In Figure \ref{fig:res_before} we plot this residual as a function of raw chip coordinate.
 Each black dot represents a single star measurecl in sinele exposure., Each black dot represents a single star measured in single exposure.
 We combinei all the exposures together to see the overall treuds., We combine all the exposures together to see the overall trends.
 Iudividual residuals were then averaged within small (—113-pixel size) cells (conveniently defined as described in the next section) aud are plotted as arrow vectors (maguified by a factor of 2500)., Individual residuals were then averaged within small $\sim$ 113-pixel size) cells (conveniently defined as described in the next section) and are plotted as arrow vectors (magnified by a factor of 2500).
 These residuals exhibit a pattern with unexpected abrupt discontinuities (deuoted with the solid. red lines)., These residuals exhibit a pattern with unexpected abrupt discontinuities (denoted with the solid red lines).
 We note that while some of the large-scale trends could be partially removed by adopting higher-order polynomials. these cliscoutinuities. which have about the same amplitude as the remaining large-scale treuds (70.02 pixels). would still remain.," We note that while some of the large-scale trends could be partially removed by adopting higher-order polynomials, these discontinuities, which have about the same amplitude as the remaining large-scale trends $\sim$ 0.02 pixels), would still remain."
 These residuals are very similar to those seen in lIXozhurina-Platais et ((2010)., These residuals are very similar to those seen in Kozhurina-Platais et (2010).
 ]t is interesting to note that these ~0.02 pixel systematic treuds are perfectly consistent. witli the lareer-than-expectec residual dispersion already. uoted in Paper LI. These fine-scale trends were simply washed out by the large internal motions of the w Ceu stars (“0.15 WEC3/UVIS pixel) over the 7-year baseline between the reference-[rame observations of GO-9L12 in 2002 aud the WEC3/UVIS observations of PID-11152 in 2009., It is interesting to note that these $\sim$ 0.02 pixel systematic trends are perfectly consistent with the larger-than-expected residual dispersion already noted in Paper I. These fine-scale trends were simply washed out by the large internal motions of the $\omega$ Cen stars $\sim$ 0.15 WFC3/UVIS pixel) over the 7-year baseline between the reference-frame observations of GO-9442 in 2002 and the WFC3/UVIS observations of PID-11452 in 2009.
 The top panel of Fig., The top panel of Fig.
" 2. shows the residuals à: plotted against the μι coordinates. for stars measured within a 100-pixel-tall strip. centered at fg=550. extracted fromchip[1]""."," \ref{fig:fetta} shows the residuals $\delta x$ plotted against the $x_{\rm raw}$ coordinates, for stars measured within a 100-pixel-tall strip, centered at $y_{\rm raw}=550$, extracted from."
. C'oustraining the &;c residuals within a vertical strip highlights the sharp changes in tle residual trend., Constraining the $\delta x$ residuals within a vertical strip highlights the sharp changes in the residual trend.
 Lt is important to note that these discontinuities are present in axes of the WEC3/UCVIS detectors aud not only alone a single axis. as was the case lor WEPC2 (Anderson Ixing 1999) or ACS/WEC (Anderson 2002).," It is important to note that these discontinuities are present in axes of the WFC3/UVIS detectors and not only along a single axis, as was the case for WFPC2 (Anderson King 1999) or ACS/WFC (Anderson 2002)."
 The bottom panel of Fig., The bottom panel of Fig.
 2. shows that these boundaries show up as single-row excursions of, \ref{fig:fetta} shows that these boundaries show up as single-row excursions of
"One major characteristic of the showerglass effect introduced in the helioseismic holography technique is that the local ingression and egression control measurements have larger phase shifts in magnetic regions, and also exhibit asymmetric phase shifts (Lindsey&Braun2005a,b).","One major characteristic of the showerglass effect introduced in the helioseismic holography technique is that the local ingression and egression control measurements have larger phase shifts in magnetic regions, and also exhibit asymmetric phase shifts \citep{lin05a, lin05b}."
". Our time-distance measurements (Figure 5)) have shown that, for the short annulus radius, acoustic waves inside active regions have a longer travel time, an order of 0.3 minutes, than the quiet Sun, for both the outgoing and ingoing travel times."," Our time-distance measurements (Figure \ref{fg5}) ) have shown that, for the short annulus radius, acoustic waves inside active regions have a longer travel time, an order of 0.3 minutes, than the quiet Sun, for both the outgoing and ingoing travel times."
" While for the longer annulus radius, the outgoing travel time inside active regions can be up to 1.0 minute shorter than in the quiet region, and the ingoing travel time is only about 0.3 minutes shorter."," While for the longer annulus radius, the outgoing travel time inside active regions can be up to 1.0 minute shorter than in the quiet region, and the ingoing travel time is only about 0.3 minutes shorter."
 The asymmetry in outgoing and ingoing travel times for both cases is obvious., The asymmetry in outgoing and ingoing travel times for both cases is obvious.
" However, all these numbers are sienificantly smaller than the phase shifts measured by the helioseismic holography technique (Lindsey&Braun20052).. though the differences mày come from different frequency bands and very different annulus (pupil) sizes."," However, all these numbers are significantly smaller than the phase shifts measured by the helioseismic holography technique \citep{lin05a}, though the differences may come from different frequency bands and very different annulus (pupil) sizes."
 It was suggested that acoustic signals inside active regions were phase shifted by the magnetic field (Lindsey&Braun2005a)., It was suggested that acoustic signals inside active regions were phase shifted by the magnetic field \citep{lin05a}.
". However, it is demonstrated from both Figures 5. and 8,, that for different annulus radii, the acoustic travel times inside active regions measured from time-distance technique vary significantly, from —1.0 to 0.3 minutes."," However, it is demonstrated from both Figures \ref{fg5} and \ref{fg8}, that for different annulus radii, the acoustic travel times inside active regions measured from time-distance technique vary significantly, from $-1.0$ to $0.3$ minutes."
" That is. phase shifts vary greatly depending on the annulus radii used to make measurements, or in other words, phase shifts depend largely on the depth, although one would expect that phase shift should remain relatively constant or varies little if a significant amount of phase shift is caused by the presence of surface magnetism rather than by the sunspot interior structures and dynamics."," That is, phase shifts vary greatly depending on the annulus radii used to make measurements, or in other words, phase shifts depend largely on the depth, although one would expect that phase shift should remain relatively constant or varies little if a significant amount of phase shift is caused by the presence of surface magnetism rather than by the sunspot interior structures and dynamics."
" Therefore, it is reasonable to believe that a large fraction of travel time anomalies measured by the time-distance method inside active regions are caused by real subphotospheric structures and dynamics below sunspots rather than only by the artifacts of sunspot surface magnetism."," Therefore, it is reasonable to believe that a large fraction of travel time anomalies measured by the time-distance method inside active regions are caused by real subphotospheric structures and dynamics below sunspots rather than only by the artifacts of sunspot surface magnetism."
" Certainly. the showerglass effect itself. how it may affect various types of helioseismic measurements, and how to account for this effect when it is significant are worth of further studies. including observational, theoretical and numerical studies."," Certainly, the showerglass effect itself, how it may affect various types of helioseismic measurements, and how to account for this effect when it is significant are worth of further studies, including observational, theoretical and numerical studies."
 Our experiment of measuring the second-skip acoustic travel times by annulus-annuls correlations helps to prove that the observed acoustic signals inside active regions provide useful information and do not corrupt time-distance results., Our experiment of measuring the second-skip acoustic travel times by annulus-annuls cross-correlations helps to prove that the observed acoustic signals inside active regions provide useful information and do not corrupt time-distance results.
" The mean travel times through sunspots region, computed by averaging the outgoing and ingoing travel times, are expected to be equal to half of the double-skip travel times, 75/2. because they actually measure the same signals along same ray paths passing through the sunspot interior."," The mean travel times through sunspots region, computed by averaging the outgoing and ingoing travel times, are expected to be equal to half of the double-skip travel times, $\tau_2/2$, because they actually measure the same signals along same ray paths passing through the sunspot interior."
" Previous measurements by Braun(1997) found that 7/2 was closer to the ingoing travel time than to the mean travel time, which led"," Previous measurements by \citet{bra97}
 found that $\tau_2/2$ was closer to the ingoing travel time than to the mean travel time, which led"
The 880 um emission in the II Zw 40 centre is dominated by free-free emission.,The 880 $\micron$ emission in the II Zw 40 centre is dominated by free-free emission.
" At this wavelength, the emission is usually dominated by dust on a galactic scale."," At this wavelength, the emission is usually dominated by dust on a galactic scale."
" Hereafter, we use the term *FIR luminosity"" for the dust emission luminosity integrated overum."," Hereafter, we use the term “FIR luminosity” for the dust emission luminosity integrated over."
" In 22reffig:radiopir, , weshowtherelationbetweenthemonochromaticluminbsity 1277] E 1.4GH2z(Li.4cuz) and the FIR luminosity Lrrn for the central star-forming region in Π Zw 40."," In \\ref{fig:radio_fir}, we show the relation between the monochromatic luminosity at 1.4 GHz $L_\mathrm{1.4\,GHz}$ ) and the FIR luminosity $L_\mathrm{FIR}$ for the central star-forming region in II Zw 40."
" The radio luminosity at 1.4 GHz is adoptedfrom the best-fit model for ;=3Myr (Figure 3cc), while the FIR luminosity is evaluated for Raust=100 and 250 pc (Section 5.2;; ir))."," The radio luminosity at 1.4 GHz is adoptedfrom the best-fit model for $t=3~\mathrm{Myr}$ (Figure \ref{fig:sed_age}c c), while the FIR luminosity is evaluated for $R_\mathrm{dust}=100$ and 250 pc (Section \ref{subsec:comparison}; ; \\ref{fig:sed_fir}) )."
"Sinceallthedif ferencebetweenthemodel fluxandtheobserved880um f luxisassumedtocome free—fromdust, theF IBlumingsiticsestimatgdE ""n romáhemodqsshoqltbetakenasupperlimit"," Since all the difference between the model free--free flux and the observed 880 $\micron$ flux is assumed to come from dust, the FIR luminosities estimated from the models should be taken as upper limits."
" For comparison, we also plot the observational data of BCDs in T"," For comparison, we also plot the observational data of BCDs in \\ref{fig:radio_fir}."
 hesampleistakenfrom Hopkinsetal.(2002) fortheF I Rand1.4GH zluminositiesof theentiresystem(i.e. ggloballuminosities)., The sample is taken from \citet{hopkins02} for the FIR and 1.4 GHz luminosities of the entire system global luminosities).
M GHz reffig:radio(Condonrir...etal.1998) and by the at A=60 um (large open squares).," We only adopt the objects with detections by the NRAO VLA Sky Survey (NVSS) at $\nu =1.4$ GHz \citep{condon98}
 and by the at $\lambda =60~\mu$ m (large open squares)."
" We also adopt the radio flux measured by the Faint Images of the Radio Sky at Twenty cm (FIRST;Becker,White,Helfand1995) (small open squares)."," We also adopt the radio flux measured by the Faint Images of the Radio Sky at Twenty cm \citep[FIRST;][]{becker95}
 (small open squares)."
 The FIRST flux is systematically smaller than the NVSS flux., The FIRST flux is systematically smaller than the NVSS flux.
 The probable reason which Hopkinsetal.(2002) suggest is the difference in their sensitivity to extended emission: FIRST images tend to miss extended emission whose angular size is larger than about 2 arcmin., The probable reason which \citet{hopkins02} suggest is the difference in their sensitivity to extended emission: FIRST images tend to miss extended emission whose angular size is larger than about 2 arcmin.
 The FIR luminosity of the BCD sample is estimated based on the 60 µπι and 100 jum fluxes obtained from NASA/IPAC Extragalactic Database (NED)., The FIR luminosity of the BCD sample is estimated based on the 60 $\mu$ m and 100 $\mu$ m fluxes obtained from NASA/IPAC Extragalactic Database (NED).
" If a sample BCD is not detected at 100jum, we utilize the upper limit at 100µπῃ to estimate an upper limit of Lryin."," If a sample BCD is not detected at $100~\micron$, we utilize the upper limit at $100~\micron$ to estimate an upper limit of $L_{\rm FIR}$."
 The data with upper limits are shown by crosses in Figure (L1.4GHz is taken from the NVSS data).," The data with upper limits are shown by crosses in Figure \ref{fig:radio_fir}
 $L_{\rm 1.4~GHz}$ is taken from the NVSS data)."
" We also show the global luminosities of II Zw 40, whose 1.4 GHz and FIR data are taken from Jaffeetal.(1978) and Vaderetal.(1993), respectively."," We also show the global luminosities of II Zw 40, whose 1.4 GHz and FIR data are taken from \citet{jaffe78}
 and \citet{vader93}, respectively."
" In order to quantitatively discuss the possible deviation from the standard radio—FIR relation for BCDs, we adopt the radio-to-FIR ratio as usually used (e.g.Condon1992)."," In order to quantitatively discuss the possible deviation from the standard radio–FIR relation for BCDs, we adopt the radio-to-FIR ratio as usually used \citep[e.g.][]{condon92}."
". Here we define qi1.4 and qis as The average and the standard deviation of q1.4 are calculated for the sample detected by both (60 and 100 um) and NVSS, while those of qis are determined for the sample in Huntetal.(2005)."," Here we define $q_{1.4}$ and $q_{15}$ as The average and the standard deviation of $q_{1.4}$ are calculated for the sample detected by both (60 and 100 $\micron$ ) and NVSS, while those of $q_{15}$ are determined for the sample in \citet{hunt05}."
". The averages of qi. and φις are 2.39 and 2.89, respectively, and the standard deviations are 0.19 and 0.26, respectively."," The averages of $q_{1.4}$ and $q_{15}$ are 2.39 and 2.89, respectively, and the standard deviations are 0.19 and 0.26, respectively."
 We also show the lines with q1.4=2.39+0.19 and q15=2.89£0.26 to show how much the radio—FIR relation in the central part of II Zw 40 is deviated from the reference radio-FIR relation for the BCDs., We also show the lines with $q_{1.4}=2.39\pm 0.19$ and $q_{15}=2.89\pm 0.26$ to show how much the radio–FIR relation in the central part of II Zw 40 is deviated from the reference radio–FIR relation for the BCDs.
" From reffig:radiosir, , weobservethatthecentralregiono f II Zw40hasasuppressed1.4 —FIRrelationofglobalBCDs, iftheFI RluminosityofII Zw40hasavaluenea "," From \\ref{fig:radio_fir}, , we observe that the central region of II Zw 40 has a suppressed 1.4 GHz luminosity and is located below the radio--FIR relation of global BCDs, if the FIR luminosity of II Zw 40 has a value near to the upper limit."
freeabsorptioninourmodels(Fig., This is interpreted to be due to free--free absorption in our models \\ref{fig:sed_age}) ).
" To avoid free-free absorption, we use the 15 GHz luminosity."," To avoid free–free absorption, we use the 15 GHz luminosity."
 The best-fit model for £=3 Myr is used for the 15 GHz luminosity of the II Zw 40 centre (Section 5.2))., The best-fit model for $t=3$ Myr is used for the 15 GHz luminosity of the II Zw 40 centre (Section \ref{subsec:comparison})).
 This luminosity fits the observationaldata at 15 GHz by Becketal. (2002).., This luminosity fits the observationaldata at 15 GHz by \citet{beck02}. .
" For comparison, we show the global luminosities of BCDs compiled"," For comparison, we show the global luminosities of BCDs compiled"
of the source spectral emission properties at lower energies (we relied on XRT data for the energy range kkeV).,of the source spectral emission properties at lower energies (we relied on XRT data for the energy range keV).
" In the considered energy range, the PCA spectrum could be modeled by a power-law with I'1.5—2.0 (see Fig. 3,,"," In the considered energy range, the PCA spectrum could be modeled by a power-law with $\Gamma$$\sim$ 1.5–2.0 (see Fig. \ref{fig:spectra},"
 upper panels)., upper panels).
" The PL photon index increased ((~1.7-1.9) during the earlier phases of the outburst MMJD, blue diamonds), became steeper (('~2.0) around the peak of the outburst MMJD, black dots), decreased until MMJD, and remained roughly constant afterwards."," The PL photon index increased $\Gamma$$\simeq$ 1.7-1.9) during the earlier phases of the outburst MJD, blue diamonds), became steeper $\Gamma$$\simeq$ 2.0) around the peak of the outburst MJD, black dots), decreased until MJD, and remained roughly constant afterwards."
" The flux of this PL component was suppressed at the outburst peak (black dots), increased during the spectral hardening until 55 MMJD (green stars), and finally decreases steadily (red triangles)."," The flux of this PL component was suppressed at the outburst peak (black dots), increased during the spectral hardening until $\sim$ MJD (green stars), and finally decreases steadily (red triangles)."
" These results agree fairly well with those obtained from aand The timing analysis was performed using the events in the energy range kkeV. Power spectra (PSD) were calculated using during intervals of ss, averaged together, and geometrically rebinned."," These results agree fairly well with those obtained from and The timing analysis was performed using the events in the energy range keV. Power spectra (PSD) were calculated using during intervals of s, averaged together, and geometrically rebinned."
" In preliminary runs, we used 2? ss time bins and verified that above 30 HHz the signal was always consistent with white noise."," In preliminary runs, we used $2^{-10}$ s time bins and verified that above $\sim$ Hz the signal was always consistent with white noise."
" The latter, also corrected for the dead-time contribution, was subtracted from the final PSD, which we computed from the background-corrected light curves binned at 2~° ss. We averaged the PSDs in long intervals corresponding to the different phases of the outburst and verified that the PSD of the single observations within these intervals presented similar characteristics."," The latter, also corrected for the dead-time contribution, was subtracted from the final PSD, which we computed from the background-corrected light curves binned at $2^{-6}$ s. We averaged the PSDs in long intervals corresponding to the different phases of the outburst and verified that the PSD of the single observations within these intervals presented similar characteristics."
" As shown in Fig. 4,,"," As shown in Fig. \ref{fig:psd},"
" all PSDs are characterized by a flat-top, band- noise that can be described using Lorentzian curves (Fig. 4))."," all PSDs are characterized by a flat-top, band-limited noise that can be described using Lorentzian curves (Fig. \ref{fig:psd}) )."
 The PSD is dominated by two broad components whose centroid frequencies move coherently throughout the outburst in rough correlation with the source flux., The PSD is dominated by two broad components whose centroid frequencies move coherently throughout the outburst in rough correlation with the source flux.
" In the early phase of the outburst 8814 MJD) the lower component peaks at ~0.5 HHz, while at the maximum of the outburst MMJD) it moves to ~1 HHz."," In the early phase of the outburst 814 MJD) the lower component peaks at $\sim$ Hz, while at the maximum of the outburst MJD) it moves to $\sim$ Hz."
" We note that the higher frequency component splits at the outburst peak into narrower features centered at 5.0+0.1 HHz and 8.4+0.6 HHz, with quality factors of 5+1 and 3+1, respectively (fractional RMS of ~8%,, uncertainties at 16)."," We note that the higher frequency component splits at the outburst peak into narrower features centered at $5.0\pm0.1$ Hz and $8.4\pm0.6$ Hz, with quality factors of $5\pm1$ and $3\pm1$, respectively (fractional RMS of $\sim8$, uncertainties at $1\sigma$ )."
" Integrating the PSD on shorter intervals, we were unable to verify wether these features are produced by narrower peaks moving in frequency, owing to the limited statistics."," Integrating the PSD on shorter intervals, we were unable to verify wether these features are produced by narrower peaks moving in frequency, owing to the limited statistics."
" At the later phases, we averaged the PSD over the intervals 8833 MJD and 8848 MJD to study possible variations of the PSD, to show a progressive decrease of the noise peak frequency (we neglected the following observations because of the lower S/N)."," At the later phases, we averaged the PSD over the intervals 833 MJD and 848 MJD to study possible variations of the PSD, to show a progressive decrease of the noise peak frequency (we neglected the following observations because of the lower S/N)."
" In the lower panels of Fig. 3,,"," In the lower panels of Fig. \ref{fig:spectra},"
" we show the integrated fractional RMS computed from the background-subtracted light curves using bins of 2-*ss as (Σο:-5»?XX,o2)(QNT) (i is the rate in the i-th bin with error σι, N=2048 is the number of bins, and r is the mean count rate)."," we show the integrated fractional RMS computed from the background-subtracted light curves using bins of $2^{-4}$ s as $\left(\sum_{i=0}^{N}(r_i-\bar r)^2 - \sum_{i=0}^{N} \sigma_i^2\right)^{1/2} / (N \bar r)$ $r_i$ is the rate in the $i$ -th bin with error $\sigma_i$, $N=2048$ is the number of bins, and $\bar r$ is the mean count rate)."
 The RMS of each observation and the corresponding uncertainty was estimated from the average and standard deviation of the values determined in each ss long time intervals., The RMS of each observation and the corresponding uncertainty was estimated from the average and standard deviation of the values determined in each s long time intervals.
" The fractional RMS remains above at the early and late stages of the outburst, while it is slightly suppressed (~18%)) at the peak (black dots)?.."," The fractional RMS remains above at the early and late stages of the outburst, while it is slightly suppressed $\sim$ ) at the peak (black ."
"Even when the ""grevo scattering.e we find a cliflerence of the temperature structures with/without scattering if we take the physical heigh as the coordinate as shown in Figure 3.","Even when the “grey” scattering, we find a difference of the temperature structures with/without scattering if we take the physical height as the coordinate as shown in Figure 3."
 The diamonds in Figure 32 indicate the lower boundary of the super-heat«(d laver. at which the Planck mean extinction optical dexh with the stellar effective temperature (7;=3.000 I) is equal to the grazing anele.," The diamonds in Figure 3 indicate the lower boundary of the super-heated layer, at which the Planck mean extinction optical depth with the stellar effective temperature $T_*=3,000$ K) is equal to the grazing angle."
 We find that with scattering. the height of the super-heated laver is always enhanced: scatering causes more Iaring disc (seealsoDullemone&Natta2003b).," We find that with scattering, the height of the super-heated layer is always enhanced; scattering causes more flaring disc \citep[see also][]{dul03}."
. Phis suggests that the scattering may allect the eloxil structure of the disc. which will be discussed in a future work.," This suggests that the scattering may affect the global structure of the disc, which will be discussed in a future work."
 1n order to understand the numerical results presented in the previous section. we here develop an analytic model as an extension of the seminal two-laver model by COG97: threc-laver model with scattering.," In order to understand the numerical results presented in the previous section, we here develop an analytic model as an extension of the seminal two-layer model by CG97: three-layer model with scattering."
 To describe Iluxes across the boundaries of the lavers. we adopt a two-stream Iddington approximation with isotropic scattering.," To describe fluxes across the boundaries of the layers, we adopt a two-stream Eddington approximation with isotropic scattering."
 The notations in this section are summarised in Table 1., The notations in this section are summarised in Table 1.
 Suppose two or three lavers in an annulus as shown in Figure 4., Suppose two or three layers in an annulus as shown in Figure 4.
 We assume that each [aver is isothermal with the temperature determined by the radiation equilibrium in the laver., We assume that each layer is isothermal with the temperature determined by the radiation equilibrium in the layer.
 Then. we consider that each laver emits the radiation characterised by its temperature and other lavers just. work as absorption and scattering media for the radiation.," Then, we consider that each layer emits the radiation characterised by its temperature and other layers just work as absorption and scattering media for the radiation."
 The radiation [rom the central star is characterised. by the stellar effective temperature., The radiation from the central star is characterised by the stellar effective temperature.
 Phe characterisation of the radiation is done by the Planck mean and the characteristic frequencies are. denoted. by cach subseript such as 77 for the stellar radiation (see the caption of Figure 4 and ‘Table 1)., The characterisation of the radiation is done by the Planck mean and the characteristic frequencies are denoted by each subscript such as “*” for the stellar radiation (see the caption of Figure 4 and Table 1).
 We denote. for example. the extinction. Cross section characterised by the stellar cllective temperature as ST.," We denote, for example, the extinction cross section characterised by the stellar effective temperature as $\kappa_*^{\rm ext}$."
 Note that all the quantities depending on the [requency are Planck averaged in this section., Note that all the quantities depending on the frequency are Planck averaged in this section.
 We call the two or three lavers. super-heatec Iaver. middle laver. ancl interior as shown in Figure 4.," We call the two or three layers super-heated layer, middle layer, and interior as shown in Figure 4."
 Phe super-heated [aver is defined. as only the laver exposed. by. the direct. stellar radiation entering into the annulus with a small erazing angle a., The super-heated layer is defined as only the layer exposed by the direct stellar radiation entering into the annulus with a small grazing angle $\alpha$ .
 The optical thickness of the [aver is =à às shown by the numerical solutions in 2., The optical thickness of the layer is $\approx\alpha$ as shown by the numerical solutions in 2.2.
" “Thus. we define the thickness of the super-heated laver as 72%""=STUNM.= a. where M. ds the gas column density. of the aver."," Thus, we define the thickness of the super-heated layer as $\tau_{\rm s,*}^{\rm ext}\equiv\kappa_*^{\rm ext}\Sigma_{\rm s}=\alpha$ , where $\Sigma_{\rm s}$ is the gas column density of the layer."
 The interior represents the isothermal part. [ound in he numerical solutions., The interior represents the isothermal part found in the numerical solutions.
 Thus. the boundary can be defined w the photosphere of its own radiation.," Thus, the boundary can be defined by the photosphere of its own radiation."
 However. we here simply. define the interior as the part other than the reated and the middle lavers.," However, we here simply define the interior as the part other than the super-heated and the middle layers."
 The middle layer is introduced. bv the following consideration., The middle layer is introduced by the following consideration.
 When the opacity cocllicient decreases. as he wavelength increases. the absorption of the radiation rom the warm super-heatecl laver. occurs well above the »hotosphere of the cold interior radiation.," When the opacity coefficient decreases as the wavelength increases, the absorption of the radiation from the warm super-heated layer occurs well above the photosphere of the cold interior radiation."
" In this case. the interior is not warmed directly by the super-heated [aver oit by the ""middle"" laver where the radiation of the super-reateck [laver is. effectively absorbed."," In this case, the interior is not warmed directly by the super-heated layer but by the “middle” layer where the radiation of the super-heated layer is effectively absorbed."
" We here define. the hickness of the middle laver as τος—sYM=1 with he eas column density of the middle laver X, although this definition is rather arbitrary."," We here define the thickness of the middle layer as $\tau_{\rm m,s}^{\rm ext}\equiv\kappa_{\rm s}^{\rm ext}\Sigma_{\rm m}=1$ with the gas column density of the middle layer $\Sigma_{\rm m}$ although this definition is rather arbitrary."
 On the other hand. when the opacity coellicient is grey. the middle laver with the above hickness is optically thick for its own radiation.," On the other hand, when the opacity coefficient is grey, the middle layer with the above thickness is optically thick for its own radiation."
 Fhus. the micelle lawer reaches the thermal equilibrium and is merged into the isothermal interior.," Thus, the middle layer reaches the thermal equilibrium and is merged into the isothermal interior."
 In. this case.we do not. need o consider the middle laver.," In this case,we do not need to consider the middle layer."
" Pherefore. we have two cases: he three-laver model when 7554,ext—iοκMa,ένο€1 and the wo-laver model when zaextc1 (or mS;ext=QUNMyecl xeause the middle laver is merged. into the interior)."," Therefore, we have two cases: the three-layer model when $\tau_{\rm m,m}^{\rm ext}\equiv\kappa_{\rm m}^{\rm ext}\Sigma_{\rm m}<1$ and the two-layer model when $\tau_{\rm m,m}^{\rm ext}\simeq1$ (or $\tau_{\rm m,i}^{\rm ext}\equiv\kappa_{\rm i}^{\rm ext}\Sigma_{\rm m}\simeq1$ because the middle layer is merged into the interior)."
 In other words. three lavers are needed when 8:fasext«Land wo layers are sullicient when 5;/87c1.," In other words, three layers are needed when $\kappa_{\rm m}^{\rm ext}/\kappa_{\rm s}^{\rm ext}<1$ and two layers are sufficient when $\kappa_{\rm i}^{\rm ext}/\kappa_{\rm s}^{\rm ext}\simeq1$."
 Lables 2 and 3 are sumniaries of the thickness of the upper two Lavers and he condition of the two or three-Iaver models., Tables 2 and 3 are summaries of the thickness of the upper two layers and the condition of the two or three-layer models.
 Importantly. he seminal two-laver model is valid only when the opacity coellicient is erev. in the frequencies interest.," Importantly, the seminal two-layer model is valid only when the opacity coefficient is grey in the frequencies interest."
 This fact has not seemed to be known well so far., This fact has not seemed to be known well so far.
" When the grazing angle à is small. the stellar lux at the top of the annulus is with the integrated. Planck function B.=(eupíz)12 and the cilution factor M,=©,/4e. where ©, is the solid angle of the stellarphotosphere from the top of the annulus."," When the grazing angle $\alpha$ is small, the stellar flux at the top of the annulus is with the integrated Planck function $B_*=(\sigma_{\rm SB}/\pi)T_*^4$ and the dilution factor $W_*=\Omega_*/4\pi$, where $\Omega_*$ is the solid angle of the stellarphotosphere from the top of the annulus."
" Η onlv the fraction fi; of the stellar photosphere is visible cause of the optically thick disc. the solid angle becomes O,=fastRy? where A, is the stellar radius and £2 is he radius of the annulus."," If only the fraction $f_{\rm vis}$ of the stellar photosphere is visible because of the optically thick disc, the solid angle becomes $\Omega_*=f_{\rm vis}\pi(R_*/R)^2$, where $R_*$ is the stellar radius and $R$ is the radius of the annulus."
 When there is scattering. a part of the incident. stellar lux is reflectecl upwarels and downwards by the superheated aver.," When there is scattering, a part of the incident stellar flux is reflected upwards and downwards by the super-heated layer."
 Calvetetal.(1991). presented an analytic expression of the scattered (lux for isotropic scattering.," \cite{cal91}
 presented an analytic expression of the scattered flux for isotropic scattering."
 From equation (5) in Calvetetal.(1991).. the outbound Huxes at the upper ancl lower boundaries of the super-heated laver (see Figure 4) 2CCONLC and where ως is the single scattering albedo at the stellar pequency and ὧν=Ylwy.," From equation (5) in \cite{cal91}, the outbound fluxes at the upper and lower boundaries of the super-heated layer (see Figure 4) become and where $\omega_*$ is the single scattering albedo at the stellar frequency and $\chi_*=\sqrt{1-\omega_*}$."
 In the derivation of equations (6) and (7). wehave assumed 7X!—ow!=a «ll. acloptecl a different upper boundary. condition from Calvet line... and neglected the term e.1 for the clownwards Dux.," In the derivation of equations (6) and (7), wehave assumed $\tau_{\rm s,*}^{\rm ext}\equiv\kappa_*^{\rm ext}\Sigma_{\rm s}=\alpha\ll1$ adopted a different upper boundary condition from \cite{cal91} , and neglected the term $e^{-1}$ for the downwards flux."
" Note that the otal scattered flux is 447""|1/72analyBy=wu? and LL=fps0 if us0 (no scattering case)"," Note that the total scattered flux is $H^{\rm up}_* + H^{\rm down}_* 
= \omega_* \alpha W_* B_* = \omega_* H^{\rm in}_*$ and $H_*^{\rm up}=H_*^{\rm down}=0$ if $\omega_*=0$ (no scattering case)."
2001).. but it is helpful to reiterate them here in one unified. picture.,", but it is helpful to reiterate them here in one unified picture."
 In terms of matching. the observations. of. the ECs14026. stars and objects like 127116. this figure shows that at these high Zi. radial modes. (and. hence. p-modes. in general) are not excited. even for very high. Z (οἱ.115.39.inCharpinetetal. 2001)..," In terms of matching the observations of the EC14026 stars and objects like $-14^{\circ}116$, this figure shows that at these high $T_{\rm eff}$, radial modes (and hence p-modes in general) are not excited, even for very high $Z$ \citep[cf. Fig.3 in ][]{Cha01}. ."
 Despite the excitation of higher- modes with slightly bluer blue edges. the extension to the instability.D zone Is Dinsullicient.," Despite the excitation of higher-order modes with slightly bluer blue edges, the extension to the instability zone is insufficient."
2 This has a natural explanation., This has a natural explanation.
 In attempting to shift the instability boundary. Z has been increased. in. order to increase the driving by the Fe-opacity bump.," In attempting to shift the instability boundary, $Z$ has been increased in order to increase the driving by the Fe-opacity bump."
 At the same time the mean molecular weight of the gas ancl the mean opacity in the layers above the driving zone has been increased., At the same time the mean molecular weight of the gas and the mean opacity in the layers above the driving zone has been increased.
 Therefore. the bFe-bunip will be located correspondingly closer to the stellar surface., Therefore the Fe-bump will be located correspondingly closer to the stellar surface.
" In. order. to make the blue edge bluer. it is necessary to maintain the contribution""EN to drivingD. from⋅ iron.⊀without.. increasing. the"," In order to make the blue edge bluer, it is necessary to maintain the contribution to driving from ironwithout increasing the"
definition of q (Eq. 7)),definition of $q$ (Eq. \ref{qdefinition}) )
 are phase-independent., are phase-independent.
 The value of q is a useful characterization of ihe mode in this case. and Fig.," The value of $q$ is a useful characterization of the mode in this case, and Fig."
 3. clearly illustrates Chat one of the two solutions for given Ris a relatively weak wave. with q~lew. whereas the other is stronger. with qc107.," \ref{comparisonlincirc} clearly illustrates that one of the two solutions for given $R$ is a relatively weak wave, with $q\sim\,$ few, whereas the other is stronger, with $q\sim10^{-2}$."
 The properties of linearly polarized waves which carry a non-zero phnase-averaged inagnelic field in the toroidal direction are shown in Fig. 4.., The properties of linearly polarized waves which carry a non-zero phase-averaged magnetic field in the toroidal direction are shown in Fig. \ref{lineardispersion}.
 The refractive index is plotted for superluminal waves Chat correspond (o ji—10100. a=100. and four different values of the parameter 5. that describes the relative amount of phase-averaged magnetic Εις carried in (he subhuninal wave (see Eq. 25)).," The refractive index is plotted for superluminal waves that correspond to $\mu=10100$ , $\sigma=100$, and four different values of the parameter $b$, that describes the relative amount of phase-averaged magnetic flux carried in the subluminal wave (see Eq. \ref{bdefinition}) )."
 It can be seen that the cutoff moves to greater /? as 5 rises., It can be seen that the cutoff moves to greater $R$ as $b$ rises.
 In the striped wind model. (his corresponds to moving away ΠΟΠ (he equator. where b—0.," In the striped wind model, this corresponds to moving away from the equator, where $b=0$."
 Although the refractive indices for non-zero b have similar shape to those for b=0. there is one potentially important difference: for b>0.65 the lower branch cuts across the ὃς=0 axis. and continues towards solutions with negative phase speed.," Although the refractive indices for non-zero $b$ have similar shape to those for $b=0$, there is one potentially important difference: for $b>0.65$ the lower branch cuts across the $\beta_>=0$ axis, and continues towards solutions with negative phase speed."
 Zero refractive index indicates that the H-DIramne coincides with the f[frame., Zero refractive index indicates that the H-frame coincides with the frame.
 These solutions are standing waves in (he pulsar wind. with constant magnetic [field ancl wavelength 2»ric.," These solutions are standing waves in the pulsar wind, with constant magnetic field and wavelength $\gg r_{\rm LC}$."
 Negative relractive index indicates propagation towards (the pulsar. in the sense that the velocity of (he LI-Draae is inward propagating.," Negative refractive index indicates propagation towards the pulsar, in the sense that the velocity of the H-frame is inward propagating."
 Nevertheless these modes still carry the particle. energy. momentum and magnetic fluxes outwards towards the nebula.," Nevertheless these modes still carry the particle, energy, momentum and magnetic fluxes outwards towards the nebula."
 Conversion of a subluminal wave into a superluminal wave is accompanied by a substantial transfer of enerev from the fields to the particles., Conversion of a subluminal wave into a superluminal wave is accompanied by a substantial transfer of energy from the fields to the particles.
 This is quantified in Fie. 5..," This is quantified in Fig. \ref{gammalin},"
 where the phase-averaged Lorentz factor (5) seen in the [frame is plotted as a function of RH. for the same values of ji. o. and b as in Fig. 4..," where the phase-averaged Lorentz factor $\langle \gamma'\rangle$ seen in the frame is plotted as a function of $R$, for the same values of $\mu$, $\sigma$, and $b$ as in Fig. \ref{lineardispersion}."
 For the smaller 6 values. b=0 and b=0.2. the upper branch of the solution has (5)2ji which implies that almost all of the field energy is converted to particlesin the equatorial regions of the wind.," For the smaller $b$ values, $b=0$ and $b=0.2$, the upper branch of the solution has $\langle \gamma'
\rangle \approx \mu$, which implies that almost all of the field energy is converted to particlesin the equatorial regions of the wind."
 At higher latitudes. only a part of the Povunting flux can be converted. since the phase-averaged," At higher latitudes, only a part of the Poynting flux can be converted, since the phase-averaged"
point spread function is not spatially invariant as is implicit in that method.,point spread function is not spatially invariant as is implicit in that method.
 It either requires the solution of a matrix (measurement) equation 2010).. which is computationally more demanding than traditional self-calibration or. as we will show. it can be written as a specific three dimensional Fourier transform.," It either requires the solution of a matrix (measurement) equation , which is computationally more demanding than traditional self-calibration or, as we will show, it can be written as a specific three dimensional Fourier transform."
 The main goal of this paper. however. is not to introduce a specific solution scheme or algorithm to these equations. but to provide more useful guidance to future methods of self-consistent three dimensional modeling of the tonosphere and its use in calibration.," The main goal of this paper, however, is not to introduce a specific solution scheme or algorithm to these equations, but to provide more useful guidance to future methods of self-consistent three dimensional modeling of the ionosphere and its use in calibration."
 The second goal is to deseribe several effects of a three dimensional tonosphere on interferometric measurements at low frequencies that go beyond thin phase-screen models. which have thus far been very successful. but are demonstratively incorrect for an extended three dimensional 1onosphere.," The second goal is to describe several effects of a three dimensional ionosphere on interferometric measurements at low frequencies that go beyond thin phase-screen models, which have thus far been very successful, but are demonstratively incorrect for an extended three dimensional ionosphere."
 The outline of the paper is as IIn Section 2. the theory of scattering of a plane-wave electric field Is restated in terms more familiar to radio interferometry and extended to the case of multiple point sources.," The outline of the paper is as In Section 2, the theory of scattering of a plane-wave electric field is restated in terms more familiar to radio interferometry and extended to the case of multiple point sources."
 In Section 3. the cross-correlation of the scattered electric field is determined tthe visibility function) and several effects of the ionosphere on imaging are described sspeckle).," In Section 3, the cross-correlation of the scattered electric field is determined the visibility function) and several effects of the ionosphere on imaging are described speckle)."
 In Section 4. the first order effect of the thickness of the ionosphere is analyzed.," In Section 4, the first order effect of the thickness of the ionosphere is analyzed."
 In Section 5. the description is extended from multiple points sources to a continuous intensity field and a mathematical expression is derived that allows one to build à three dimensional model of the power spectrum of the tonosphere from information obtained only in the interferometer plane aa holographic principle).," In Section 5, the description is extended from multiple points sources to a continuous intensity field and a mathematical expression is derived that allows one to build a three dimensional model of the power spectrum of the ionosphere from information obtained only in the interferometer plane a holographic principle)."
 In Section 6. if is shown how the tomographic method. connects to an extension of the phase-sereen approach to three dimensions using a Radon transformation and applying the Fourier projection-slice theorem.," In Section 6, it is shown how the tomographic method, connects to an extension of the phase-screen approach to three dimensions using a Radon transformation and applying the Fourier projection-slice theorem."
 In Section 7. we summarize our results and give conclusions.," In Section 7, we summarize our results and give conclusions."
 In this section. the basic theory of electric-field tomography of a weakly scattering media ts restated as first introduced by1969).," In this section, the basic theory of electric-field tomography of a weakly scattering media is restated as first introduced by."
. The notation is adapted to that typically used in radio astronomy and radio interferometry2001)., The notation is adapted to that typically used in radio astronomy and radio interferometry.
". Under the assumption that the refractive index of the 10nosphere varies slowly over a single photon wavelength. and the tonosphere is ""frozen"" over the time scale the radiation passes through it. one can write the electric field equation in a dielectric medium as where 2,(r) is the refractive index of the medium at a position r and at frequency v. and &=27/.t is the usual wave number for a wavelength j|2c/v1999).."," Under the assumption that the refractive index of the ionosphere varies slowly over a single photon wavelength, and the ionosphere is “frozen” over the time scale the radiation passes through it, one can write the electric field equation in a dielectric medium as where $n_{\nu}(\vc{r})$ is the refractive index of the medium at a position ${\vc r}$ and at frequency $\nu$, and $k=2\pi/\lambda$ is the usual wave number for a wavelength $\lambda = c/\nu$."
 In a plasma with electron density Πω) the refractive index is given by a=1[nie[ρε]1-Qv. where is the plasma frequency. being typically -5 MMHz for the ionosphere.," In a plasma with electron density $n_{e}(\vc{r})$ the refractive index is given by $n^{2}={1-[n_{e} e^{2}/(\nu^{2} m_{e} \epsilon_{0})]}\equiv {1-(\nu_{\rm p}/\nu)^{2}}$, where $ \nu_{\rm p}$ is the plasma frequency, being typically $\sim$ MHz for the ionosphere."
vp We drop the explicit frequency dependence of the electric field. but it should implicitly be assumed.," We drop the explicit frequency dependence of the electric field, but it should implicitly be assumed."
 Because the three components of the electric field are independent. each component satisfies the same solution and we can simply use the scalar £ to describe the electric field.," Because the three components of the electric field are independent, each component satisfies the same solution and we can simply use the scalar $E$ to describe the electric field."
 If the refractive index of the ronosphere is near unity. which is often the case in radio astronomy v>> the equation for the electric field can be conveniently rewrittenvy). as where P(r)=£pi-1]/4z is called the p," If the refractive index of the ionosphere is near unity, which is often the case in radio astronomy $\nu \gg \nu_{\rm p}$ ), the equation for the electric field can be conveniently rewritten as where $\Phi(\vc{r})\equiv k^{2}[n^{2} - 1]/4 \pi$ is called the ."
otential. In case of 1«I. the scattering potential strength is close to zero and scattering isweak?.," In case of $n\approx1$, the scattering potential strength is close to zero and scattering is."
. When @(r)=0 everywhere. the equation reduces to the Helmholtz equation for a plane wave.," When $\Phi(\vc{r}) = 0$ everywhere, the equation reduces to the Helmholtz equation for a plane wave."
" The solution of this equation can be obtained through Green's functions and leads to an implicit integral equation where E(r)=ΕΟΟ+E'""(r).", The solution of this equation can be obtained through Green's functions and leads to an implicit integral equation where $E(\vc{r}) = E^{(i)}(\vc{r}) + E^{(s)}(\vc{r})$.
 The first term is the incident (plane) wave and the second term the scattered wave., The first term is the incident (plane) wave and the second term the scattered wave.
 The latter is equal to the integral term above and carried out over the entire volume V in which 5.#|., The latter is equal to the integral term above and carried out over the entire volume $V$ in which $n\neq 1$.
" The last factor in the integrand indicates that E""(r) is a spherically outgoing wave. assuming the extent of the scattering potential is small compared to the distance between observer and scattering medium."," The last factor in the integrand indicates that $E^{(s)}(\vc{r})$ is a spherically outgoing wave, assuming the extent of the scattering potential is small compared to the distance between observer and scattering medium."
 We now extend the single plane-wave description of to an incident electric field that results from the sum of N point sources that satisfy the solution of the Helmholtz equation in free space., We now extend the single plane-wave description of to an incident electric field that results from the sum of $N$ point sources that satisfy the solution of the Helmholtz equation in free space.
 Later in the paper we further extend this to a continuous intensity field., Later in the paper we further extend this to a continuous intensity field.
 We also express all physical distances in units of 2. 1e. u=rit(uv) in the remainder of the paper.," We also express all physical distances in units of $\lambda$, i.e. $\vc{u} \equiv \vc{r}/\lambda = (u,v,w)$ in the remainder of the paper."
 Its Fourier equivalent is s(Sy.5j. 549.," Its Fourier equivalent is $\vc{s}=(s_{u},s_{v},s_{w})$ ."
" The incident electric field in this case becomes where S,, is the ux-density of point source n=1...N and so, are unit vectors that point in the directions of the point sources."," The incident electric field in this case becomes where $S_{n}$ is the flux-density of point source $n=1\dots N$ and $\vc{s}_{0,n}$ are unit vectors that point in the directions of the point sources."
 These point sources are (for now) assumed to dominate the electric. field and can be compared to thecalibrators in radio interferometry., These point sources are (for now) assumed to dominate the electric field and can be compared to the in radio interferometry.
 In the weak scattering approximation. the scattered wave has a relatively low amplitude compared to the incident wave.," In the weak scattering approximation, the scattered wave has a relatively low amplitude compared to the incident wave."
 We can then replace E(u) with E'(u) to first-order (Born) approximation. finding where the subscript indicates that the scattered field is that to first order Born approximation and d(u)I=[ri—1].," We can then replace $E(\vc{u})$ with $E^{(i)}(\vc{u})$ to first-order (Born) approximation, finding where the subscript indicates that the scattered field is that to first order Born approximation and $\Phi(\vc{u}) \equiv [n^{2}-1]$."
 We note here that the second order can be obtained by substituting Eu+EY'(u) secondin to the original integral equation to obtain a solution to order. etc.," We note here that the second order can be obtained by substituting $E^{(i)}(\vc{u}) + E_{1}^{(s)}(\vc{u})$ in to the original integral equation to obtain a solution to second order, etc."
 This iterative scheme only works for weak scattering.where the values of @(u) do not exceed unity.," This iterative scheme only works for weak scattering,where the values of $\Phi(\vc{u})$ do not exceed unity."
 If we use the fact that the spherical outgoing wave can be written as, If we use the fact that the spherical outgoing wave can be written as
AGN activity 3. 4). 2 normal. galaxy candidates exhibiting narrow emission line optical spectra (έτη. 8) and. 2 unclassified sources with no optical spectroscopic information 110. 11).,"AGN activity 3, 4), 2 `normal' galaxy candidates exhibiting narrow emission line optical spectra 7, 8) and 2 unclassified sources with no optical spectroscopic information 10, 11)."
 The X-ray/optical properties of this latter class of sources suggests ACN activity., The X-ray/optical properties of this latter class of sources suggests AGN activity.
" We also note that although the X-ray and. optical properties of ""normal galaxy candidates are consistent with stellar origin for the X-ray emission we cannot exclude the possibility of heavily obscured AGN or LLAGN.", We also note that although the X-ray and optical properties of `normal' galaxy candidates are consistent with stellar origin for the X-ray emission we cannot exclude the possibility of heavily obscured AGN or LLAGN.
 Summarising the classification above. 9 out of the 12 X-ravradio matches are associated with AGN activity on the basis of the broad optical emission lines (total of 3). radio morphology (2) or X-rayfoptical properties (4).," Summarising the classification above, 9 out of the 12 X-ray/radio matches are associated with AGN activity on the basis of the broad optical emission lines (total of 3), radio morphology (2) or X-ray/optical properties (4)."
" One radio source is associated with X-rav cluster emission with the remaining two sources being ""normal galaxy canclidates.", One radio source is associated with X-ray cluster emission with the remaining two sources being `normal' galaxy candidates.
 Brinkmann et al. (, Brinkmann et al. (
2000) cross-correlated the FIRST radio survey with the RASS-LL catalogue witha kkeV limiting flux of =10Mvestem7.,"2000) cross-correlated the FIRST radio survey with the RASS-II catalogue with a keV limiting flux of $\approx10^{-13}\,\rm erg\,s^{-1}\,cm^{-2}$."
 These authors found, These authors found
and D4;,and $B_{abs}$.
 Differcuces between the predicted redshifts from the two equations are all much smaller than the scatter in either., Differences between the predicted redshifts from the two equations are all much smaller than the scatter in either.
 Another possibility is that the external πουτι around SCRBs plavs some role in their Iuuünositv., Another possibility is that the external medium around SGRBs plays some role in their luminosity.
 The interstellar mediuni is typically less dense iu elliptical ealaxies than in actively star forming galaxies., The interstellar medium is typically less dense in elliptical galaxies than in actively star forming galaxies.
 However. this should ouly affect the observed fux from external shocks that GRB ejecta drive into the ambient imediuim.," However, this should only affect the observed flux from external shocks that GRB ejecta drive into the ambient medium."
 External shocks are believed. to power the afterglow cluission. but not the GRD Inuinosity itself (e.¢.. Piran 2005).," External shocks are believed to power the afterglow emission, but not the GRB luminosity itself (e.g., Piran 2005)."
 Sinee our observed auticorrelation iuvolves tli SCRB οταν flux itself. we disfavor this explanation.," Since our observed anticorrelation involves the SGRB $\gamma$ -ray flux itself, we disfavor this explanation."
 To conclude. we have found au auticorrelation between the energy of short ganna ray burst and the Iuuinosity of its host ealaxy.," To conclude, we have found an anticorrelation between the energy of a short gamma ray burst and the luminosity of its host galaxy."
a Such an anticorrelation occurs witoe a probability of ~1% in simulations that account fo observational selection effects. aud so is probably real.," Such an anticorrelation occurs with a probability of $\sim 1\%$ in simulations that account for observational selection effects, and so is probably real."
 Its physical origin is unclear. though it is most likely dic to a correlation between the age of au SCRD progenitor and the luminosity of the explosion.," Its physical origin is unclear, though it is most likely due to a correlation between the age of an SGRB progenitor and the luminosity of the explosion."
 Hf this correlation is à contiuuous distribution. if can provide approximate redshift estimates for SGRDs.," If this correlation is a continuous distribution, it can provide approximate redshift estimates for SGRBs."
 Time will tell whether this correlation holds for a larger sample. and whether it is a property of a single distribution. or if it instead reflects an underlying division of SCRBs into two plivsicallv distinct sets.," Time will tell whether this correlation holds for a larger sample, and whether it is a property of a single distribution, or if it instead reflects an underlying division of SGRBs into two physically distinct sets."
 Ithauk Saugecta Malhotra. Evan Scaunapieco. Patrick Young. Sununer Starfield. aud Frauk Tiunes for stimulating discussions: Saudra Savaglio aud Edo Berger for their published host ealaxy compilations: R. Quin. E. MeMalion. aud J. Murphy for the CRBlog database at http://erad[0.as.utexas.edu/erblos.php: aud the Max Planck Institute for Astronomy for hospitality chiving completion of this work.," I thank Sangeeta Malhotra, Evan Scannapieco, Patrick Young, Sumner Starrfield, and Frank Timmes for stimulating discussions; Sandra Savaglio and Edo Berger for their published host galaxy compilations; R. Quimby, E. McMahon, and J. Murphy for the GRBlog database at http://grad40.as.utexas.edu/grblog.php; and the Max Planck Institute for Astronomy for hospitality during completion of this work."
the value E(VI)=0.03+0.02. which will be the assumed reddening throughout this paper.,"the value $E(V-I)=0.03\pm0.02$, which will be the assumed reddening throughout this paper."
 Distance moduli of the Palomar class clusters have been often overestimated in the past., Distance moduli of the Palomar class clusters have been often overestimated in the past.
 Inani Rosine (1962)) searched Palomar 12 for variables., Kinman Rosino \cite{k62}) ) searched Palomar 12 for variables.
 They found three variables. oue of them previously discovered by Zwickv (1957)).," They found three variables, one of them previously discovered by Zwicky \cite{z57}) )."
 Based on the mean apparent magnitude of these RR Lyrae. Pal 12 was initially located farther than 50 kpe from the Calactic ceuter (Iburis 1976)).," Based on the mean apparent magnitude of these RR Lyrae, Pal 12 was initially located farther than 50 kpc from the Galactic center (Harris \cite{h76}) )."
 It is only after IICSO photometric study that a nore precise distance modulus has been eiven (about 114 kpce). on the basis of the V iiaguitude of the poorly populated ΠΟ.," It is only after HC80 photometric study that a more precise distance modulus has been given (about 14 kpc), on the basis of the $V$ magnitude of the poorly populated HB."
 We derive the distance to Pal 12 by comparing its IB with that of NOC 6362. which is the only GC at ΤΟΠ)= with measured a-clemeuts abundance (cf.," We derive the distance to Pal 12 by comparing its HB with that of NGC 6362, which is the only GC at ${\rm [Fe/H]}\simeq -1$ with measured $\alpha$ -elements abundance (cf."
 Tab., Tab.
 2 in Carney. 1996).," 2 in Carney, 1996)."
 Piotto et al. (1998)), Piotto et al. \cite{p98}) )
 eive MV(IB)=0.7840.05 for NGC 6362: this value isπο represcutative of the Pal 12 IIB huuinosity. since we must correct for the age (cf," give $M_V({\rm HB})=\mvhba \pm \errmva$ for NGC 6362; this value is representative of the Pal 12 HB luminosity, since we must correct for the age (cf."
 Sect. 1) , Sect. \ref{age}) )
and à abundance offsets between both clusters., and $\alpha$ abundance offsets between both clusters.
 A decrease in age mÓuplies an increase in the IID stars nias and hDnuuainositv. the exact dependency beiue a function of Z.," A decrease in age implies an increase in the HB stars mass and luminosity, the exact dependency being a function of $Z$ ."
 Although no Z=0.002 (the Pal 12 metallicity) models are available. au interpolation from the Z=0.001 and Z=0101. Bertelli ct al. 1991. ," Although no $Z=0.002$ (the Pal 12 metallicity) models are available, an interpolation from the $Z=0.001$ and $Z=0.004$, Bertelli et al. \cite{b94}, ,"
hereafter. BOL) model isochrones leads to estimate a chauge AAA~0.07 mag. which reduces the age by AVA∙≻ (cf," hereafter B94) model isochrones leads to estimate a change $\Delta M_V \sim -0.07$ mag, which reduces the age by $30\%$ (cf."
 Sect. 1)., Sect. \ref{age}) ).
 Spectroscopy of 2 NCC 6362 eiauts has been obtained by Cratton (1987)). who measured loFo=0.32+0.09.," Spectroscopy of 2 NGC 6362 giants has been obtained by Gratton \cite{g87}) ), who measured ${\rm [\alpha/Fe]}=0.32\pm0.09$."
 Tn view of the results by Brown ct al. (1997)), In view of the results by Brown et al. \cite{b97}) )
" presented in Sec. δεν,"," presented in Sec. \ref{metallicity},"
" a comparison of the Pal 12 CAID with NGC 6362 niust take iuto account the ""o-enlaucemeut"" of the latter.", a comparison of the Pal 12 CMD with NGC 6362 must take into account the $\alpha$ -enhancement” of the latter.
 As iscussed m more detail iu Sect. L.," As discussed in more detail in Sect. \ref{age},"
 an imercase of 0.3 dex iu [a/Fe] mimics an increase of 0.2 dex in the equivalent [Fe/II]. aud implies a decrease in he IIB brightness.," an increase of 0.3 dex in ${\rm [\alpha/Fe]}$ mimics an increase of 0.2 dex in the equivalent [Fe/H], and implies a decrease in the HB brightness."
 The exact value depeuds ou the slope of the huninositv-netallicity relation for the IB., The exact value depends on the slope of the luminosity-metallicity relation for the HB.
 Although this is still controversial. a typical value AAS/AlFe/T]=0.20 can be used (Carney et al. 1992)).," Although this is still controversial, a typical value $\Delta
M_V/\Delta {\rm [Fe/H]}=0.20$ can be used (Carney et al. \cite{c92}) ),"
 which therefore meas AAA=0.01 mag in our case., which therefore means $\Delta M_V=0.04$ mag in our case.
 We should also take iuto account possible differences in the mass loss rates along the RGB between the two clusters., We should also take into account possible differences in the mass loss rates along the RGB between the two clusters.
 These would affect the ZATIB mass. aud. hence its luminosity.," These would affect the ZAHB mass, and hence its luminosity."
 Iu order to coustrain such an effect. we cau compare the colors of the red IID of Pal 12 and NCC 6362.," In order to constrain such an effect, we can compare the colors of the red HB of Pal 12 and NGC 6362."
 Tudeed. using again the D91 isochrones we find that. in the red HB region. a change in theZATIB mass of OLAS. will change the WB location of a star bv |0.22 mae in (B. V)aud O.OF mag in V.," Indeed, using again the B94 isochrones we find that, in the red HB region, a change in theZAHB mass of $M_\odot$ will change the HB location of a star by $+0.22$ mag in $(B-V)$ and $-0.07$ mag in $V$ ."
 The effect is therefore three times larecr in the (0T) color than in the V. magnitude., The effect is therefore three times larger in the $(B-V)$ color than in the $V$ magnitude.
 The actual dereddened colors of the red TBs of the two clusters ave (D.Vg~0.73 for Pal 12 (Stetson et al.," The actual dereddened colors of the red HBs of the two clusters are $(B-V)_0
\sim 0.73$ for Pal 12 (Stetson et al."
 1989). and (CD.V)g~0.51 for NGC 6362 (Piotto et al.," 1989), and $(B-V)_0 \sim 0.54$ for NGC 6362 (Piotto et al."
 1995)., 1998).
 IIence. a color difference of ~0.2 agin (BV) exists between Pal 12 aud NGC 6362. which corresponds toa <OLAS. mass loss difference.," Hence, a color difference of $\sim
0.2$ mag in $(B-V)$ exists between Pal 12 and NGC 6362, which corresponds to a $<0.1 M_\odot$ mass loss difference."
 Towever. this higher IID amass for Pal 12 is consistent with its lower age.," However, this higher HB mass for Pal 12 is consistent with its lower age."
 According to BOL. the turnoff mass of a cluster will change by 0.LAZ.. if its age is changed wo 5 Cor.," According to B94, the turnoff mass of a cluster will change by $\sim 0.1 M_\odot$ if its age is changed by $\sim 5$ Gyr."
 Since. in the BOL scale. the typical GC age would be f~L1:15 Gyr (Saviane et al. 1998)).," Since, in the B94 scale, the typical GC age would be $t\sim 14\div15$ Gyr (Saviane et al. \cite{savetal98}) ),"
 the higher uass of the Pal 12 IID is casily explaiue« by its ~30% ower age (c£., the higher mass of the Pal 12 HB is easily explained by its $\sim 30\%$ lower age (cf.
 Sect. 11., Sect. \ref{age}) ).
 A mass loss differcutial correction is therefore not needed., A mass loss differential correction is therefore not needed.
" Iu sunuuary. we expect that the Pal 12 UB should v0 0,07 nae brighter than that of NGC 6362 in view of it» vounger age and 0.04 mae brighter due to its lower a clement content. 1.0. AA(IIB)=0.67£0.05, "," In summary, we expect that the Pal 12 HB should be 0.07 mag brighter than that of NGC 6362 in view of its younger age and 0.04 mag brighter due to its lower $\alpha$ element content, i.e. $M_V({\rm HB})=
\mvhbb \pm \errmvb$ ."
As the appareu magnitude of the Pal 12 IIB is Vip=lr.l8£0.02 (where the error has been computed takine into accouir the calibration uncertainties). the apparent distance modulus becomesnAA=16.51.," As the apparent magnitude of the Pal 12 HB is $V_{\rm HB} = \vhb \pm
\errvhb$ (where the error has been computed taking into account the calibration uncertainties), the apparent distance modulus becomes$m-M_V=\apparent$."
 Caven the asstuned reddening £L(BVj=0.02. the absolute distance uocdulhus is ηAMjy=16.15d0.10.," Given the assumed reddening $E(B-V)=0.02$, the absolute distance modulus is $(m-M)_0=\mmzero \pm \errmm$."
 The estimate of the error icludes the uncertainties ou the calibration zero-))mt. on the magnitude of the NGC 6362 IID. aud on he absorption.," The estimate of the error includes the uncertainties on the calibration zero-point, on the magnitude of the NGC 6362 HB, and on the absorption."
 Our value of the distance to Pal 12 is xrfectlv compatible with previous estimates: IICSO. eive 16.22:0.35. GOSS 16.1:16.5. S89 16.3. and DADO 16.16 or the absolute distance modulus.," Our value of the distance to Pal 12 is perfectly compatible with previous estimates: HC80 give $16.2\pm0.35$, GO88 $16.1\div16.5$, S89 $16.3$, and DA90 $16.46$ for the absolute distance modulus."
 The adopted distance modulus corresponds to a distance from the Sun HR.=19.5d:0.9 IxXpc. a distance roni the Galactic center Ree=16.240.7 kpc. aud a icieli5 Zop=τι.05 below he Calactic plane (ve adopted a distance from the Sun to the Galactic center %=S.0L05 kpe: Reid 1993)).," The adopted distance modulus corresponds to a distance from the Sun $R_\odot=\ds\pm\eds$ Kpc, a distance from the Galactic center $R_{GC}=\dgc\pm\edgc$ kpc, and a height $Z_{GP}=\dgp\pm\edgp$ below the Galactic plane (we adopted a distance from the Sun to the Galactic center $R=8.0\pm0.5$ kpc; Reid \cite{r93}) )."
 The first evidence of a relatively voung age for Pal 12 was given by GOss. who estimated that Pal 12 must be vounecr than 17 Tuc ou the basis of an atypically siuall value of the magnitude difference between the TB aud the TO.," The first evidence of a relatively young age for Pal 12 was given by GO88, who estimated that Pal 12 must be younger than 47 Tuc on the basis of an atypically small value of the magnitude difference between the HB and the TO."
 Indeed. this was the first clear identification of a vouug GGC.," Indeed, this was the first clear identification of a young GGC."
 AhnDnost at the une time. S89 presented. an indepeudeut BV CCD study.," Almost at the same time, S89 presented an independent $BV$ CCD study."
 They compared the Pal 12 fiducial ROB o those of T Tue and M5. which bracket Pal 12 inctallicity. concluding that no match could be oua.," They compared the Pal 12 fiducial RGB to those of 47 Tuc and M5, which bracket Pal 12 metallicity, concluding that no match could be found."
 The simplestexplanation was that Pal 12 is vounger than he other two clusters by some , The simplestexplanation was that Pal 12 is younger than the other two clusters by some .
Iu bothstudies. the 17 Tuc fiducial lines were taken youn Hesser e al. (1987.. ," In bothstudies, the 47 Tuc fiducial lines were taken from Hesser et al. \cite{h87}, ,"
hereafter Ηδη., hereafter H87).
 These fiducials were constructed by mereing 2 aud V.CCD photometiy or SSOO stars below the MS turnoff to the evolved part of heCMD conmiug from earlier photographic work (κα, These fiducials were constructed by merging $B$ and $V$CCD photometry for 8800 stars below the MS turnoff to the evolved part of theCMD coming from earlier photographic work (Hesser
closer to the Dohnanyi(1969) distribution.,closer to the \citet{dohn69} distribution.
 The average slope of the size distributions for the collisional models is —3.5., The average slope of the size distributions for the collisional models is $-3.5$.
 The curves for the collisional models also contain dips in the mid-sized grain populations., The curves for the collisional models also contain dips in the mid-sized grain populations.
 The dips become deeper as the optical depth increases., The dips become deeper as the optical depth increases.
" Stark&Kuchner(2008) hypothesized that in the absence of a resonant source population, grains with size such their PR time, tpa, matched their collision time, ἔσομ, would dominate the optical depth of a resonant ring."," \citet{star08} hypothesized that in the absence of a resonant source population, grains with size such their PR time, $t_{\rm PR}$, matched their collision time, $t_{\rm coll}$, would dominate the optical depth of a resonant ring."
 Our new simulations seem to confirm this hypothesis., Our new simulations seem to confirm this hypothesis.
" We will call this critical grain size s,.", We will call this critical grain size $s_c$.
" Table 1 shows s, in each simulation, averaged over the first 50 records in each stream."," Table \ref{tab:criticalgrainsize} shows $s_{c}$ in each simulation, averaged over the first 50 records in each stream."
 The critical grain size is also plotted in Figure 6 for each simulation as a circle on the corresponding black curve., The critical grain size is also plotted in Figure \ref{fig:grainsizes} for each simulation as a circle on the corresponding black curve.
with: Even i£ wiclely used. it has been recently shown that the AlAlerritt parameterization does not fit well the anisotropy parameter. measured in. cosmological bbody simulations (Diaferio1999:Colinctal.2000:Rasiaetal.2004:Diemandct 2004).,": Even if widely used, it has been recently shown that the Merritt parameterization does not fit well the anisotropy parameter measured in cosmological body simulations \cite{Dia99,Col00,RTM,Die04}."
. Indeed. Alamon Lokas (2004a.b) have shown that the following paramoeterization works:: with rug—runfr. and rare a characteristic anisotropy radius.," Indeed, Mamon Lokas (2004a,b) have shown that the following parameterization works with $x_{ML} = r_{ML}/r_s$ and $r_{ML}$ a characteristic anisotropy radius."
 With the Mamon Lokas parameterization of the anisotropy. parameter. the radial velocity. dispersion. turns out tobe: Actually. it is worth noting that some of the simulations considered by Mamon Lokas are tailored on cluster scale haloes (Rasiaetal.2004) or on the subhaloes (Diemandetal.2004) so that it is at least. premature to reject the AMlerritt model for galaxy scale svstems.," With the Mamon Lokas parameterization of the anisotropy parameter, the radial velocity dispersion turns out to: Actually, it is worth noting that some of the simulations considered by Mamon Lokas are tailored on cluster scale haloes \cite{RTM} or on the subhaloes \cite{Die04} so that it is at least premature to reject the Merritt model for galaxy scale systems."
" Hence. in reffig: anisigma.. we report the racial velocity dispersion 0, for the NOt model for the isotropic. model. the MMerrit/ ancl the Manion Lokas anisotropic parameterizations."," Hence, in \\ref{fig: anisigma}, we report the radial velocity dispersion $\sigma_r$ for the N04 model for the isotropic model, the Merrit and the Mamon Lokas anisotropic parameterizations."
 As expected. anisotropic mocdels predict an higher velocity dispersion. with the AlAlerritt scheme being more elective in raising σε.," As expected, anisotropic models predict an higher velocity dispersion with the Merritt scheme being more effective in raising $\sigma_r$."
 In. particular. it turns out that the lower is ro. the higher is the the radial velocity dispersion that may be qualitatively explained. by noting that the model becomes more and. more isotropic as roa inereases.," In particular, it turns out that the lower is $x_{OM}$, the higher is the the radial velocity dispersion that may be qualitatively explained by noting that the model becomes more and more isotropic as $x_{OM}$ increases."
 We have checked that these results do not depend on 5., We have checked that these results do not depend on $\gamma$.
 Note that there is a degeneracy between mass and anisotropy so that it is dillicult to choose among isotropic. Αλλο and Mamon Lokas mocels on the basis of the radial velocity dispersion only.," Note that there is a degeneracy between mass and anisotropy so that it is difficult to choose among isotropic, Merritt and Mamon Lokas models on the basis of the radial velocity dispersion only."
 Moreover. the differences among these models may be partially washed out when considering the projected velocity dispersion since this enters also other parameters (related to the Luminous components) thus leading to possible degeneracies that seriously complicate the interpretation of the data.," Moreover, the differences among these models may be partially washed out when considering the projected velocity dispersion since this enters also other parameters (related to the luminous components) thus leading to possible degeneracies that seriously complicate the interpretation of the data."
 Introclucing an anisotropy in the velocity space also changes the DIE ofthe model that now becomes a function not only of the binding energv €. but also of the mocdulus L of the angular momentum. vector.," Introducing an anisotropy in the velocity space also changes the DF of the model that now becomes a function not only of the binding energy ${\cal{E}}$, but also of the modulus $L$ of the angular momentum vector."
 The construction of general anisotropic DE f(£.22L) for spherically svmmetric dedensity pairs is discussed in detail in Dejonghe (1986).," The construction of general anisotropic DF $f({\cal{E}}, L)$ for spherically symmetric density pairs is discussed in detail in Dejonghe (1986)."
 Llowever. here. we restrict our attention to the MMerrit. parameterization in which case the DE depends on & anc £ onlythrough the ::Lt is then possible to compute the DE using the following generalization of the Eddington formula with," However, here, we restrict our attention to the Merrit parameterization in which case the DF depends on ${\cal{E}}$ and $L$ onlythrough the :It is then possible to compute the DF using the following generalization of the Eddington formula \cite{D86,BT87} : with"
Therefore the only source of magnetic field generation is the first term on the right hand side of Ίσα. (AL4)).,Therefore the only source of magnetic field generation is the first term on the right hand side of Eq. \ref{mag_gen_eq}) ).
 Eq. CX14)), Eq. \ref{mag_gen_eq}) )
 can therefore be simplified to: This equation can be used to ect an order-of-magnitude estimate of the generated. magnetic field., can therefore be simplified to: This equation can be used to get an order-of-magnitude estimate of the generated magnetic field.
 Eq. CX12)), Eq. \ref{tcoupl}) )
 can be used to calculate. to the zeroth order. the evolution equation of electron. velocity Ποιά (Pechbles Yu 1970).," can be used to calculate, to the zeroth order, the evolution equation of electron velocity field (Peebles Yu 1970)."
 This equation alone with the continuity equation. (leq. (5))), This equation along with the continuity equation (Eq. \ref{cont_eq}) ))
 and VE=0. Eq. (BG6)).," and ${\bf \nabla .E} = 0$, Eq. \ref{diver_e}) ),"
" ancl dropping all terms proportional to D. gives: This equation can be solved along with the evolution equation of 6, by WIND approximation and these solutions can be matched to large scale solutions (lu Sugivama 1995)."," and dropping all terms proportional to ${\bf B}$, gives: This equation can be solved along with the evolution equation of $\delta_\gamma$ by WKB approximation and these solutions can be matched to large scale solutions (Hu Sugiyama 1995)."
 We discuss here approximate solutions at cilferent epochs., We discuss here approximate solutions at different epochs.
 First we discuss solutions during phase (ο) of the evolution., First we discuss solutions during phase (c) of the evolution.
 We note that all the terms on the right hand side except for the 9» term are smaller as compared to this term for scales smaller than the sound horizon scale c1/31)., We note that all the terms on the right hand side except for the $\delta_\gamma$ term are smaller as compared to this term for scales smaller than the sound horizon scale $\simeq 1/\sqrt{3} \eta$.
 Phe electron pressure is always negligible as compared to the photon pressure in the pre-recombination universe., The electron pressure is always negligible as compared to the photon pressure in the pre-recombination universe.
 With these simplification and. bearing in mind that 1/72«1 during the evolution. Eq. CX16))," With these simplification and bearing in mind that $1/R \ll 1$ during the evolution, Eq. \ref{elec_evol}) )"
 is solved to give: Llere we have only retained the solution compatible with adiabatic initial conditions (see e.g. Llu Sugivama 1995) and In phase (d) of the evolution of the plasma. the tight-coupling approximation breaks down and the photon-slip which damps perturbations (Silk damping) must be taken into account.," is solved to give: Here we have only retained the solution compatible with adiabatic initial conditions (see e.g. Hu Sugiyama 1995) and In phase (d) of the evolution of the plasma, the tight-coupling approximation breaks down and the photon-slip which damps perturbations (Silk damping) must be taken into account."
 The solution including the Silk damping is (see e.g. Peebles 19801: Llere. nsPhe silk. damping. scale. at any epoch. can be obtained. from⋅ this. expression:. in the matter-dominated era.," The solution including the Silk damping is (see e.g. Peebles 1980): Here, The silk damping scale, at any epoch, can be obtained from this expression: $k_{\rm silk} \simeq (15/(2\tau_{\gamma e}\eta))^{1/2} 
\simeq (4 \, {\rm Mpc})^{-1} ((1+z)/10^3)^{-5/4}$ in the matter-dominated era."
 The velocity [eld in the linear evolution remains non-vortical. and hence can be found from the continuity equation (leq. (5))).," The velocity field in the linear evolution remains non-vortical, and hence can be found from the continuity equation (Eq. \ref{cont_eq}) ))."
 lt should. be noted that solutions for the barvon density ancl velocity fields differ from the corresponding quantities for the electrons only by replacing 2 defined here as 4=tp./(3pi) (seo c.g. Llu Suegivama 1995)., It should be noted that solutions for the baryon density and velocity fields differ from the corresponding quantities for the electrons only by replacing $R$ defined here as $R' = 4\rho_\gamma/(3\rho_b)$ (see e.g. Hu Sugiyama 1995).
 As 2x1 for the evolution of the plasma in the pre-recombination universe. for barvonic. densities compatible with primordial nucleosvnthesis. the barvon and electron quantities can be used interchangeably in evaluating the second order expression above.," As $R' \gg 1$ for the evolution of the plasma in the pre-recombination universe, for baryonic densities compatible with primordial nucleosynthesis, the baryon and electron quantities can be used interchangeably in evaluating the second order expression above."
 The evolution of electron density and velocity in the oscillatory regime and the super-horizon solutions. prior to the epoch of recombination. can be summarized as (for solutions at super-horizon scales in this conformal-Newton gauge see e.g. Ma )ortschinger 1995): Llere RD and MD correspond to radiation and matter dominated epochs. respectively.," The evolution of electron density and velocity in the oscillatory regime and the super-horizon solutions, prior to the epoch of recombination, can be summarized as (for solutions at super-horizon scales in this conformal-Newton gauge see e.g. Ma Bertschinger 1995): Here ${\rm RD}$ and ${\rm MD}$ correspond to radiation and matter dominated epochs, respectively."
We begin by excising the anomalous points near the peak and fitting the remaining light curve to the degenerate form of the profile that obtains in the high-magnification limit 1996).. Here F'(f) is the instantaneous flux. F4. is the peak flux above the baseline of the model. Fi. is the baseline flux. ty is the time of the peak. and έως 1s the effective width of the light curve.,"	We begin by excising the anomalous points near the peak and fitting the remaining light curve to the degenerate form of the profile that obtains in the high-magnification limit , Here $F(t)$ is the instantaneous flux, $F_{\rm peak}$ is the peak flux above the baseline of the model, $F_{\rm base}$ is the baseline flux, $t_0$ is the time of the peak, and $t_{\rm eff}$ is the effective width of the light curve."
 These degenerate parameters are related to the standard parameters by Fia=Pio. Fia=FY|Fy. and fae=ugfp. where fg is the Einstein timescale. i.e.. the time required for the source to move an Einstein radius.," These degenerate parameters are related to the standard parameters by $F_{\rm peak} = F_{\rm s}/u_0$, $F_{\rm base} = F_{\rm s} + F_{\rm b}$, and $t_{\rm eff} = u_0 t_{\rm E}$, where $t_{\rm E}$ is the Einstein timescale, i.e., the time required for the source to move an Einstein radius."
 We find fy~2151593.6 (Heliocentric Julian Date). £4r~1day. and that Γρ roughly corresponds to Ziaj15.3 (Fig. [1] D.," We find $t_0\sim 2451393.6$ (Heliocentric Julian Date), $t_{\rm eff}\sim 1\ \mbox{day}$, and that $F_{\rm peak}$ roughly corresponds to $I_{\rm peak} \sim 15.3$ (Fig. \ref{fig:data}] ])."
 However. the exact values of those parameters are quite dependent on which points we excise.," However, the exact values of those parameters are quite dependent on which points we excise."
" Note that though the baseline magnitude i$ measured to be 7,4.=19.1. there is at this point essentially no information of how the baseline flux divides into source flux F; and blended flux Fj,. which is why it is necessary to fit the light curve to the degenerate profile (eq. [1 )."," Note that though the baseline magnitude is measured to be $I_{\rm base} = 19.1$, there is at this point essentially no information of how the baseline flux divides into source flux $F_{\rm s}$ and blended flux $F_{\rm b}$ , which is why it is necessary to fit the light curve to the degenerate profile (eq. \ref{eqn:degen}] ])."
 Next. we reinsert the excised points near the peak and note that the maximum (over-magnified) deviation from the degenerate fitis ~0.5mag (Fig. 1).," 	Next, we reinsert the excised points near the peak and note that the maximum (over-magnified) deviation from the degenerate fit is $\sim 0.5\ \mbox{mag}$ (Fig. \ref{fig:data}) )."
 We then establish a set of initial trial event geometries as follows., We then establish a set of initial trial event geometries as follows.
" For each diamond-shaped caustic produced by various geometries of (¢/.4) pairs. we examine the magnification as a function of distance from the “caustic center"" along each of the three directions defined by the four cusps of the caustic (one direction is redundant due to the reflection symmetry with respect to the binary axis)."," For each diamond-shaped caustic produced by various geometries of $d$ $q$ ) pairs, we examine the magnification as a function of distance from the “caustic center” along each of the three directions defined by the four cusps of the caustic (one direction is redundant due to the reflection symmetry with respect to the binary axis)."
 Here for computational simplicity. we use an analytic proxy point for the “caustic center.” defined as follows.," Here for computational simplicity, we use an analytic proxy point for the “caustic center,” defined as follows."
 For close binaries. we adopt the center of mass of the lens system for the caustic center.," For close binaries, we adopt the center of mass of the lens system for the caustic center."
 On the other hand. for wide binaries. the position of one component of the binary shifted toward the other component by |/(1|q1)|+ is chosen as the caustic center. [," On the other hand, for wide binaries, the position of one component of the binary shifted toward the other component by $[d(1+q^{-1})]^{-1}$ is chosen as the caustic center. ["
These choices can be understood from the approximation developed in Appendix Α..,These choices can be understood from the approximation developed in Appendix \ref{sec:ap}.
 See also (1999).., See also .]
| Then. at each point αν. the point of the cusp-axis crossing. we take the ratio of the actual magnification to that predicted for the corresponding PSPL lens: a point mass lens located at the caustic center with the same mass as either the whole binary (close binaries) or the nearest lens alone (wide binaries).," Then, at each point $\vect{u}_\bullet$, the point of the cusp-axis crossing, we take the ratio of the actual magnification to that predicted for the corresponding PSPL lens: a point mass lens located at the caustic center with the same mass as either the whole binary (close binaries) or the nearest lens alone (wide binaries)."
 We proceed with the search only if this ratio lies in the interval |1.2.1.5].," We proceed with the search only if this ratio lies in the interval $[1.2,1.8]$."
" We then choose a source trajectory perpendicular to the cusp axis as an initial trial model. with wy=[u,| and te=fof[u,|. and use four-parameter (o. £g. fy. and «9) downhill simplex to search for a best-fit model to minimize (7. which is determined by a linear-flux fit to the model of lens magnifications (and thus the corresponding F. and Fi, are found immediately). subject to the constraint of fixed (7.4)."," We then choose a source trajectory perpendicular to the cusp axis as an initial trial model, with $u_0 = |\vect{u}_\bullet|$ and $t_{\rm E} = t_{\rm eff}/|\vect{u}_\bullet|$, and use four-parameter $\alpha$, $t_{\rm E}$, $t_0$ , and $u_0$ ) downhill simplex to search for a best-fit model to minimize $\chi^2$, which is determined by a linear-flux fit to the model of lens magnifications (and thus the corresponding $F_{\rm s}$ and $F_{\rm b}$ are found immediately), subject to the constraint of fixed $d$ $q$ )."
 Figure 2. shows the result of this search (based on ISIS solutions) in a contour plot of \7-surface over the (d.q)-space., Figure \ref{fig:chi2cont} shows the result of this search (based on ISIS solutions) in a contour plot of $\chi^2$ -surface over the $d$ $q$ )-space.
 The 4? is rising around all the boundaries shown in Figure 2.. except toward lower q.," The $\chi^2$ is rising around all the boundaries shown in Figure \ref{fig:chi2cont}, except toward lower $q$ ."
 While the \?-surface flattens with Ay?232 as the mass ratio ¢ becomes very small (q.<0.01) for d~0.065 and d—15. we find no evidence of a decrease of the 4? às 4 becomes smaller than 0.01.," While the $\chi^2$ -surface flattens with $\Delta\chi^2\simeq 32$ as the mass ratio $q$ becomes very small $q\la 0.01$ ) for $d\sim 0.065$ and $d\sim15$, we find no evidence of a decrease of the $\chi^2$ as $q$ becomes smaller than $0.01$."
 Rather it asymptotically approaches Ay?~32., Rather it asymptotically approaches $\Delta\chi^2\simeq 32$.
 We find well-localized minima of 4? over the searched (d.4)-space. one of close binaries (d.~ 0.13) and the other of wide binaries (4~ 11.3). whose exact parameters depend slightly on whether we use the ISIS or DoPHOT photometry (Tables 2. and 3)).," We find well-localized minima of $\chi^2$ over the searched $d$ $q$ )-space, one of close binaries $d\simeq 0.13$ ) and the other of wide binaries $d\simeq 11.3$ ), whose exact parameters depend slightly on whether we use the ISIS or DoPHOT photometry (Tables \ref{tab:model} and \ref{tab:modelw}) )."
 We adopt the ISIS solutions in the subsequent discussion (and they are also what is shown in Fig. 2)), We adopt the ISIS solutions in the subsequent discussion (and they are also what is shown in Fig. \ref{fig:chi2cont}) )
 because the light curve shows less scatter and consequently. the errors for the model parameter determinations are smaller.," because the light curve shows less scatter and consequently, the errors for the model parameter determinations are smaller."
 Despite the combination of the dense coverage near the peak by the PLANET data and the extensive baseline coverage by the MACHO data. the final two ISIS models (i.e.. the close binary and the wide binary) are essentially indistinguishable: A4?=0.6 for 412 degrees of freedom. with the close binary having the lower 4? of the two.," Despite the combination of the dense coverage near the peak by the PLANET data and the extensive baseline coverage by the MACHO data, the final two ISIS models (i.e., the close binary and the wide binary) are essentially indistinguishable: $\Delta\chi^2=0.6$ for 412 degrees of freedom, with the close binary having the lower $\chi^2$ of the two."
 In Appendix A.. we discuss this degeneracy in further detail.," In Appendix \ref{sec:ap}, we discuss this degeneracy in further detail."
 Finally. we also note that vote and Ενπμ from the best-fit binary-lens model parameters are not quite the same as the τη and F4; derived from the initial degenerate form of the light curve fit.," 	Finally, we also note that $u_0 t_{\rm E}$ and $F_{\rm s}/u_0$ from the best-fit binary-lens model parameters are not quite the same as the $t_{\rm eff}$ and $F_{\rm peak}$ derived from the initial degenerate form of the light curve fit."
 These discrepancies are traceable to small differences between a true curve and its degenerate approximation (eq. [1]]).," These discrepancies are traceable to small differences between a true curve and its degenerate approximation (eq. \ref{eqn:degen}] ]),"
 and are not due to differences between the binary and corresponding PSPL event. which are very small except around the peak.," and are not due to differences between the binary and corresponding PSPL event, which are very small except around the peak."
 Since the parameters derived from equation (1) function only as seeds for simplex. and since the final V ?-surface is very well-behaved. these discrepancies in inputvalues do not influence the final result.," Since the parameters derived from equation \ref{eqn:degen}) ) function only as seeds for simplex, and since the final $\chi^2$ -surface is very well-behaved, these discrepancies in inputvalues do not influence the final result."
 The mass ratios of the best fit models are y=0.310dc0.011 (close binary) and 4=0.751+0.193(wide binary). which are well away from the regime of planetary companions (¢= 0.01).By comparison. the ratio of the duration of the anomalous," 	The mass ratios of the best fit models are $q=0.340\pm 0.041$ (close binary) and $q=0.751\pm 0.193$(wide binary), which are well away from the regime of planetary companions $q\la 0.01$ ).By comparison, the ratio of the duration of the anomalous"
10 cluster has not vet. collapsed.,the cluster has not yet collapsed.
 The application of the code in its first. version. that does not. follow the detailed. temperature evolution. is thus appropriate. since less than of the gas is in regions with Jo>10! K (see also Mo White 2002).," The application of the code in its first version, that does not follow the detailed temperature evolution, is thus appropriate, since less than of the gas is in regions with $T>>10^4$ K (see also Mo White 2002)."
 In the following we present the results of our simulations., In the following we present the results of our simulations.
 In Fie., In Fig.
 4 we show illustrative slices cut. through the simulation boxes., \ref{fig04} we show illustrative slices cut through the simulation boxes.
 Phe six panels show the neutral hydrogen number density for the M3. field region (upper panels) and the 783. proto-cluster (lower panels). at redshifts. from left to right. z=16.5.12 and. 8.5.," The six panels show the neutral hydrogen number density for the `M3' field region (upper panels) and the `S3' proto-cluster (lower panels), at redshifts, from left to right, $z=16.5, 12$ and 8.5."
 The ionized regions (black) clearly grow dillerentlv. in the two environments., The ionized regions (black) clearly grow differently in the two environments.
 In the field. high density. peaks are uncommon ancl LIL regions easily. break into the IGM: in the proto-cluster the density is higher. so ionization is more cillieult and recombination is faster.," In the field, high density peaks are uncommon and HII regions easily break into the IGM; in the proto-cluster the density is higher, so ionization is more difficult and recombination is faster."
 Moreover. many photons are initially needed. to ionize the high density gas surrounding the sources. before photons can escape into the low density IGM.," Moreover, many photons are initially needed to ionize the high density gas surrounding the sources, before photons can escape into the low density IGM."
 As a result. more photons are required to ionize the proto-cluster region: although in the proto-cluster the ionizing photon production per unit mass is higher at high redshift (see Fig. 1)).," As a result, more photons are required to ionize the proto-cluster region; although in the proto-cluster the ionizing photon production per unit mass is higher at high redshift (see Fig. \ref{fig01}) ),"
 filaments of neutral gas are still present after the field region is almost completely ionized., filaments of neutral gas are still present after the field region is almost completely ionized.
 ‘This is more clearly seen in Fig. 5..," This is more clearly seen in Fig. \ref{fig05},"
" where the volume averaged ionization [raction. τρ. is shown as a function of redshift for ""M and S3."," where the volume averaged ionization fraction, $x_v$, is shown as a function of redshift for `M3' and `S3'."
" This ionization fraction is defined AS =SM,au;i where or; and V; are the ionization fraction and volume of the ;-th cell respectively. and V. is the total volume of the box."," This ionization fraction is defined as $x_v=\sum_i x_i V_i / V$, where $x_i$ and $V_i$ are the ionization fraction and volume of the $i$ -th cell respectively, and $V$ is the total volume of the box."
 The evolution proceeds cilferently in the two regions: initially the ionization fraction are comparable. but the proto-cluster gas ionizes more slowly. for the reasons discussed above.," The evolution proceeds differently in the two regions: initially the ionization fraction are comparable, but the proto-cluster gas ionizes more slowly, for the reasons discussed above."
" While the ""M3 field region already. has ar.~0.000 at 2=8.3. rs~OS for the cluster at the same recshilt."," While the `M3' field region already has $x_v \sim 0.999$ at $z=8.3$, $x_v \sim 0.8$ for the proto-cluster at the same redshift."
 A comparison between volume and. mass. averaged ionization fractions isuseful to. understand bow the reionization process proceeds., A comparison between volume and mass averaged ionization fractions isuseful to understand how the reionization process proceeds.
 We define the latter quantity aS du=»weAdAl. where M; is the mass in the ;-th cell and A is the total mass in the box.," We define the latter quantity as $x_m=\sum_i x_i M_i / M$, where $M_i$ is the mass in the $i$ -th cell and $M$ is the total mass in the box."
 In Fig., In Fig.
" ο the redshift evolution of we. (filled circles) and uy, (open circles) is compared both for the ""M. field region (upper panel) and for the ""S3. proto-cluster (lower panel).", \ref{fig06} the redshift evolution of $x_v$ (filled circles) and $x_m$ (open circles) is compared both for the `M3' field region (upper panel) and for the `S3' proto-cluster (lower panel).
" Phe most interesting feature of these plots is that. while for the field region οτι is always comparable with ri. ur,a. initially in the proto-cluster. Le. at first relatively dense cells are ionized."," The most interesting feature of these plots is that, while for the field region $x_m$ is always comparable with $x_v$, $x_m>x_v$ initially in the proto-cluster, i.e. at first relatively dense cells are ionized."
 This because the density immediately. surrounding the sources must be ionized. before the photons can break into the IGM.," This because the density immediately surrounding the sources must be ionized, before the photons can break into the IGM."
 Εις indicates that the assumption adopted by some authors (e.g. Miralda-Escudé.. Haechnelt Rees 2000). that lower density σας always ects ionized before higher density regions. fails in a cluster environment.," This indicates that the assumption adopted by some authors (e.g. Miralda-Escudé,, Haehnelt Rees 2000), that lower density gas always gets ionized before higher density regions, fails in a cluster environment."
 The trend reverses at later times. when large. low clensity regions ect ionized and the highest density peaks remain neutral or recombine.," The trend reverses at later times, when large, low density regions get ionized and the highest density peaks remain neutral or recombine."
 Although almost of the proto-cluster volume is ionized by 2~ S. this is true for only ~TO%D of the mass.," Although almost of the proto-cluster volume is ionized by $z \sim 8$ , this is true for only $\sim 70\%$ of the mass."
 Some approaches to the study of the reionization, Some approaches to the study of the reionization
considered small enough for applving linear The nonlinear cascade proceeds as it is expected in magnetized plasma (Chenetαἱ.2010)..,considered small enough for applying linear The nonlinear cascade proceeds as it is expected in magnetized plasma \citep{horbury08}.
 14 develops the well-known power anisotropy. which means that most of the energy resides. al the end of the simulation. in oblique or quasi-perpendicular modes.," It develops the well-known power anisotropy, which means that most of the energy resides, at the end of the simulation, in oblique or quasi-perpendicular modes."
 This is visible in Figure 18.. where we show the spectrum of magnetic perturbations as a contour plot in wavevector space.," This is visible in Figure \ref{fig5}, where we show the spectrum of magnetic perturbations as a contour plot in wavevector space."
 In previous works it has been argued (hat the dynamics of the dissipation at small scales can be described bv linear mechanisms. since the amplitude of fluctuations become so V.ΗΕ that non-linear effects can be neglected.," In previous works it has been argued that the dynamics of the dissipation at small scales can be described by linear mechanisms, since the amplitude of fluctuations become so small that non-linear effects can be neglected."
 It is also argued that à purely linear clamping --nechanism will result in an exponential roll-over of the spectrum of turbulent fluctuations., It is also argued that a purely linear damping mechanism will result in an exponential roll-over of the spectrum of turbulent fluctuations.
" —1 (his section. we attempt to interpret the simulation results at the light of Vlasov linear jeorv,"," In this section, we attempt to interpret the simulation results at the light of Vlasov linear theory."
 We have used a Vlasov solver to identify linear modes in the range of wavevectors of interest for this study., We have used a Vlasov solver to identify linear modes in the range of wavevectors of interest for this study.
 We anticipate that. at these scales. it becomes very diffieult to follow je numerical solution of the dispersion relation [or one single mode. as many branches of the solutions overlap.," We anticipate that, at these scales, it becomes very difficult to follow the numerical solution of the dispersion relation for one single mode, as many branches of the solutions overlap."
 We show in Figures 19-20 the dispersion relations [or whistler modes. KAW and Langmuir waves. for different angle of propagation (€=415*.605. 80*).," We show in Figures \ref{fig6}- \ref{fig7} the dispersion relations for whistler modes, KAW and Langmuir waves, for different angle of propagation $\theta=45^\circ,60^\circ,80^\circ$ )."
 We focus attention on oblique modes. since we know that the cascade tends to develop the aforementioned anisotropy.," We focus attention on oblique modes, since we know that the cascade tends to develop the aforementioned anisotropy."
 We point out that what we call Langmuir waves are the eeneralizalion of electrostatic Lanenmür modes in presence of magnetic field. studied in WillesandCairns(2000).," We point out that what we call Langmuir waves are the generalization of electrostatic Langmuir modes in presence of magnetic field, studied in \citet{willes00}."
. Thev are lightly damped hieh-frequency oscillations. which are not usually taken in account in (his The same dispersion curves for the damping rate only of KAW and whistler are plotted in log-log scale in Figure 21..," They are lightly damped high-frequency oscillations, which are not usually taken in account in this The same dispersion curves for the damping rate only of KAW and whistler are plotted in log-log scale in Figure \ref{fig8}."
 Lie£a£.(2001) have proposed an empirical fit lor the dampiug rate 5 of IKAW of the form: where ry.ie.iy are positive Πάπα parameters.," \citet{li01} have proposed an empirical fit for the damping rate $\gamma$ of KAW of the form: where $m_1,m_2,m_3$ are positive fitting parameters."
 Thev have also argued. on the basis of a simplified model. which treats the nonlinear cascade as an isotropic diffusion process in wavenumber space. (hat in order (o recover a power-law lor the spectrum of magnetic fluctuations. Che damping rate should be itself a power-law as a function of ÀÁ.," They have also argued, on the basis of a simplified model, which treats the nonlinear cascade as an isotropic diffusion process in wavenumber space, that in order to recover a power-law for the spectrum of magnetic fluctuations, the damping rate should be itself a power-law as a function of $k$."
 What is interesting. in Figure 21.. is Chat al this scale the damping rate seems to follow precisely a power law in A.," What is interesting, in Figure \ref{fig8}, is that at this scale the damping rate seems to follow precisely a power law in $k$."
 The form of Equation 4 could still be valid. but the exponential factor becomes essentially very close (o 1. as & becomes very," The form of Equation \ref{fit} could still be valid, but the exponential factor becomes essentially very close to 1, as $k$ becomes very"
 The form of Equation 4 could still be valid. but the exponential factor becomes essentially very close (o 1. as & becomes very.," The form of Equation \ref{fit} could still be valid, but the exponential factor becomes essentially very close to 1, as $k$ becomes very"
filling the observational data but are also consistent wilh the origin ancl evolution οἱ cometary dust.,fitting the observational data but are also consistent with the origin and evolution of cometary dust.
 For example. Iximuraetal.(2006.2003) considered. two tvpes of ballistic ageregales: BPCA (Ballistic Particle Cluster Ageregates) ancl BCCA (Ballistic Cluster Cluster Aggregate).," For example, \citet{kimur06, kimur03} considered two types of ballistic aggregates: BPCA (Ballistic Particle Cluster Aggregates) and BCCA (Ballistic Cluster Cluster Aggregate)."
 These two types of aggregates differ in (heir porosity., These two types of aggregates differ in their porosity.
" For an equivalent number V of constituent monomers of radius «. the BPCA aggregates are compact than DCCA. where the porosity is defined as P=1—Ααα.)5 and the characteristic radius a,=y5/3a,. where a, is the gvration radius of the aggregate (Ixozasa a,=(1/2N3)(1/2.N7-xX,tcu(ri)rU)mrGDDE.ri). with withrG) r() the position of the center of the 7i!” monomer."," For an equivalent number $N$ of constituent monomers of radius $a$, the BPCA aggregates are compact than BCCA, where the porosity is defined as $P = 1
- N(a/a_{c})^{1/3}$ and the characteristic radius $a_{c} \equiv
\sqrt{5/3}\, a_{g}$, where $a_{g}$ is the gyration radius of the aggregate \citep{kozasa92}, $a_{g} = (1/2N^{2}) \times \mathop{\sum}_{i,j=1}^{N}(r(i)-r(j))^{2}$ , with $r(i)$ the position of the center of the $i^{th}$ monomer."
 lxolokolovaetal.(2007) ancl Iximuraetal.(2006). describe the polarimetric and IR. properties of comet dust using these (wo tvpes of ballistic aggregates. consisting of sub-micron monomers (of radius e20.1 jmm)) with a IHallev-tvpe composition which includes silicates. amorphous carbon. and organic relractory material.," \citet{kol07} and \cite{kimur06} describe the polarimetric and IR properties of comet dust using these two types of ballistic aggregates, consisting of sub-micron monomers (of radius $a
\simeq 0.1$ ) with a Halley-type composition which includes silicates, amorphous carbon, and organic refractory material."
 These models can account for the general behavior of (he maxim polarization. the shape of the polarization curve as a function of phase angle P(a). and (he negative polarization branch.," These models can account for the general behavior of the maximum polarization, the shape of the polarization curve as a function of phase angle $P(\alpha)$, and the negative polarization branch."
 In addition. such models also vield the low geometric albedos (of order 25% that is tvpical for comet dust) aud red photometric colors (normalizedreflectivitygradient:Jewitt&Meech1956) that are (vpical for comet dust. for example C/2004 Q2 (Machholz) (Linetal.2007). and others discussed in Ixolokolovaetal.(2004) or Hadamcik&Levasseur-Hegourd(2000).," In addition, such models also yield the low geometric albedos (of order $\simeq 5$ that is typical for comet dust) and red photometric colors \citep[normalized reflectivity
gradient;][]{jewittmeech86} that are typical for comet dust, for example C/2004 Q2 (Machholz) \citep{lin2007} and others discussed in \citet{kolv04} or \citet{hadamcik09}."
 Indeed. Fig.," Indeed, Fig."
 2 of Joshi.Ganesh.&Balivan(2011) indicates that the colors of comet C/2007 N3 (Lulin) are not blue., 2 of \citet{joshi11} indicates that the colors of comet C/2007 N3 (Lulin) are not blue.
 For proper characterization of the observed low geometric albedos and red photometric of comets. the values of the refractive index used to describe the aggregate ensembles are crucial.," For proper characterization of the observed low geometric albedos and red photometric of comets, the values of the refractive index used to describe the aggregate ensembles are crucial."
 INimuraetal.(2006) demonstrate that if the imaginary part. &. of the index of refraction becomes smaller than 0.4 the color of comet dust becomes blue.," \citet{kimur06} demonstrate that if the imaginary part, $k$, of the index of refraction becomes smaller than 0.4 the color of comet dust becomes blue."
 The same ballistic aggregates. BPCA and.BCCA. were considered by Lasueοἱ to model comet polarization.," The same ballistic aggregates, BPCA andBCCA, were considered by \citet{lasue09} to model comet polarization."
 Mixing them with spheroidal silicate particles thev, Mixing them with spheroidal silicate particles they
within a bandpass of MALI foreground. removal therefore imposes a minimum accessible wave-number of μι20.04(1|z2)/7.5].+ 1].,within a bandpass of MHz [foreground removal therefore imposes a minimum accessible wave-number of $k_{\rm min} \approx 0.04[(1+z)/7.5]^{-1}$ $^{-1}$ ].
 We compute the sensitivity with which the effect. of quasars on the shape of the 2l-em power spectun could be detected. following the procedure outlined by ο and ? (seealso.2).., We compute the sensitivity with which the effect of quasars on the shape of the 21-cm power spectum could be detected following the procedure outlined by \cite{mcquinn2006} and \cite{bowman2006} \citep[see also ][]{wlg2008}.
 Written in terms of the cosmic wave-vector k—kpky. the resulting error in the 21-0nm power spectrum is where Ziv.sz250(1|z)niai Why is the svsteni temperature of the instrument. D(z) is the comoving distance to the point of emission at redshift z. AL is the comoving depth of the survey. volume corresponding to the bandwidth D. fy is the total integration time. (6%17) is the number density of baselines that can observe the visibility U. where €—hyD/252 and A is the observed wavelength.," Written in terms of the cosmic wave-vector $\kvec = \kvec_\parallel + \kvec_\perp$, the resulting error in the 21-cm power spectrum is where $T_{\rm sys} \approx 250[(1+z)/7]^{2.6}$ K is the system temperature of the instrument, $D(z)$ is the comoving distance to the point of emission at redshift $z$, $\Delta D$ is the comoving depth of the survey volume corresponding to the bandwidth $B$, $t_0$ is the total integration time, $n_{\rm b}(U,\nu)$ is the number density of baselines that can observe the visibility $\Uvec$, where $U = k_\perp D/2\pi$ and $\lambda$ is the observed wavelength."
 Although the observed. 21-cm power spectrum. is not spherically svimmetric. it is symmetric about the line of sight.," Although the observed 21-cm power spectrum is not spherically symmetric, it is symmetric about the line of sight."
 This makes it possible to calculate the overall power spectral sensitivity. of the radio interferometer using the Fouricr modes. contained. within an infinitesimal annulus around the line of sight of constant (4.4). where cos(8)kez/k (z is the unit vector pointing in the direction of the line of sight).," This makes it possible to calculate the overall power spectral sensitivity of the radio interferometer using the Fourier modes contained within an infinitesimal annulus around the line of sight of constant $(k,\theta)$, where $\cos(\theta) = \kvec \cdot \hat{\zvec}/k$ $\hat{\zvec}$ is the unit vector pointing in the direction of the line of sight)."
 The power spectral sensitivity over such an, The power spectral sensitivity over such an
charracteristics.,racteristics.
oed on different or tuue-varving accretion rates. we rere propose a different interpretation for the origin of he spectroscopic sub-populations of radio-lou AGN.,"based on different or time-varying accretion rates, we here propose a different interpretation for the origin of the spectroscopic sub-populations of radio-loud AGN."
 We speculate that the separation between LEG axd WEG is rot due to a different rate of accretion but. mstead. to a different of accretion.," We speculate that the separation between LEG and HEG is not due to a different rate of accretion but, instead, to a different of accretion."
 In this scenario. TEC are »owered by accretion of cold eas. ee. provided by a receut nerecr With a gas rich galaxy (6:8.Baldi&Capetti2008).," In this scenario, HEG are powered by accretion of cold gas, e.g. provided by a recent merger with a gas rich galaxy \citep[e.g. ][]{baldi08}."
". Cold eas. approaching the ceutral regions of the galaxy. ors the various structures commonly seen in these ACN, such as a molecular torus. a Broad Line Region. aud a standard. ecometrically thin. accretion disk."," Cold gas, approaching the central regions of the galaxy, forms the various structures commonly seen in these AGN, such as a molecular torus, a Broad Line Region, and a standard, geometrically thin, accretion disk."
 Conversely. LEG accrete hot material. provided by the ample reservoir of their N-rav emittiug gaseous coronae.," Conversely, LEG accrete hot material, provided by the ample reservoir of their X-ray emitting gaseous coronae."
 This process las been shown to be able to account for the nuclear activity of FR I radio-galaxies (Allenctal.2006:xDalinaverde2008) extending up to a radio power of log [τν~33.," This process has been shown to be able to account for the nuclear activity of FR I radio-galaxies \citep{allen06,balmaverde08} extending up to a radio power of log $L_{178} \sim 33$."
" The teiiperature of the accreting οas ds typically around ] keV. —10* K. This prevents the formation of the ""cold"" structures. in particular of a molecular torus. but also of the clouds of the DLR. whose ionized portion has a temperature of 10? Ix. uuless the inflowing gas is able to cool dramatically on its wav to the ceuter of the galaxy."," The temperature of the accreting gas is typically around 1 keV, $\sim 10^7$ K. This prevents the formation of the “cold” structures, in particular of a molecular torus, but also of the clouds of the BLR, whose ionized portion has a temperature of $\sim 10^5$ K, unless the in-flowing gas is able to cool dramatically on its way to the center of the galaxy."
 Iudeed DLR are not seen in LEC and they also do uot eenerallv show the hieh level of absorption iu the X-rav band expected if their uuclei were seen through au obscuring torus (6.8. Hardcastleetal. 2009))., Indeed BLR are not seen in LEG and they also do not generally show the high level of absorption in the X-ray band expected if their nuclei were seen through an obscuring torus (e.g. \citealt{hardcastle09}) ).
 Moreover. the properties of the accretion disk are ikelv to be substantially altered due to the high initial cluperature of the eas.," Moreover, the properties of the accretion disk are likely to be substantially altered due to the high initial temperature of the gas."
" From a geometrical poiut of view. a ""cold. standard accretion disk is flattened by rotation. while it remains eeometrically thick for higher enperatures,"," From a geometrical point of view, a “cold”, standard accretion disk is flattened by rotation, while it remains geometrically thick for higher temperatures."
 While the radiative cussion in a standard disk is domunatec by UV and soft X-ray photons. a rotter disk cuits most of the radiation at higher energies.," While the radiative emission in a standard disk is dominated by UV and soft X-ray photons, a hotter disk emits most of the radiation at higher energies."
 The üeher temperature iu case of hot accretion corresponds also to a lower radiative efficienev. due to the reduced eas cooling at these temperatures (but see 1992.. for the effects of high density on the gas cooling function).," The higher temperature in case of hot accretion corresponds also to a lower radiative efficiency due to the reduced gas cooling at these temperatures (but see \citealt{dumont92}, for the effects of high density on the gas cooling function)."
 Ata given accretion rate. the umber of ionizing photous is then reduced. due to the combination of higher average photon cucrey auc lower overall ciission.," At a given accretion rate, the number of ionizing photons is then reduced, due to the combination of higher average photon energy and lower overall emission."
 The cluitted ρουται is harder aud it produces lines of lower excitation. as discussed above.," The emitted spectrum is harder and it produces lines of lower excitation, as discussed above."
 This effect can produce the spectral separation between LEC and TEC sinularly to RIAF., This effect can produce the spectral separation between LEG and HEG similarly to RIAF.
 Furthermore. Capettietal.(2005) showed that the uon-thenual cores in ER I produce a sufficient flux of hie[um enerev photons to account for the ionization of their NLR.," Furthermore, \citet{capetti:cccriga} showed that the non-thermal cores in FR I produce a sufficient flux of high energy photons to account for the ionization of their NLR."
" The radio aud. optical Iuminosities of the nuclei of FR I aud ER II/LECG are linked by a common linear correlation (Chniabereeetal.2002) and. as shown by Fig. ὃν,"," The radio and optical luminosities of the nuclei of FR I and FR II/LEG are linked by a common linear correlation \citep{chiaberge:fr2} and, as shown by Fig. \ref{o3re},"
 this applies also to radio cores aud line ΕΕ, this applies also to radio cores and line luminosity.
 It is then possible that the dominant source of ionizing photous iu LEG nust be ascribed to nou-thermal emission. associated with the vase of their jets. rather then with their accretion clisks.," It is then possible that the dominant source of ionizing photons in LEG must be ascribed to non-thermal emission, associated with the base of their jets, rather then with their accretion disks."
 The accretion rate iu FR ΠΠΤα can be estimated w ooxtrapolating the scaling relation between jet and accretion power derived for low hunineositv FR I/LECG to he objects of highest huninosity of this class., The accretion rate in FR II/LEG can be estimated by extrapolating the scaling relation between jet and accretion power derived for low luminosity FR I/LEG to the objects of highest luminosity of this class.
 Baluaverde estimated that the accretion rate needed ο power a radio source with Lig~LO? cre Db ds Ῥιω1018 ore st.," \citet{balmaverde08}
 estimated that the accretion rate needed to power a radio source with $L_{178}
\sim 10^{33}$ erg $^{-1}$ is $P_{\rm accr} \sim 10^{44.6}$ erg $^{-1}$."
 This requires Pu~mi Cresyo 21 or the most powerful LEC. correspoudiug to a fraction i~0.1 of the Eddington rate for a 107NL; black hole.," This requires $P_{\rm accr} \sim 10^{46}$ erg $^{-1}$ for the most powerful LEG, corresponding to a fraction $\dot{\rm m} \sim 0.1$ of the Eddington rate for a $^9 {\rm M}_{\sun}$ black hole."
 These hot flows at high accretion rate can be probably associated with the optically thin. geometrically thick solutions (ee. Abramowiczetal.1995)) that can reach. for a Viscosity parameter a=0.1. a rate of the order of the Eddinetou onc.," These hot flows at high accretion rate can be probably associated with the optically thin, geometrically thick solutions (e.g. \citealt{abramowicz95}) ) that can reach, for a viscosity parameter $\alpha =0.1$, a rate of the order of the Eddington one."
 The niechanisni of jet launching. likely το be determined x the disk structure in the region closer to the black hole. wieght not be seusitive to the eas history. but only to the final accretion rate.," The mechanism of jet launching, likely to be determined by the disk structure in the region closer to the black hole, might not be sensitive to the gas history, but only to the final accretion rate."
 Provided that this reaches comparable high levels in WEC and LEG. racio PAructures of simular morphology aud power cau be formed oei the two sub-classes.," Provided that this reaches comparable high levels in HEG and LEG, radio structures of similar morphology and power can be formed in the two sub-classes."
 The separation between TEC aud LEG is ronuinisceut of that found bv R06 for the SDSS sources. mostly radio-quiet ACN.," The separation between HEG and LEG is reminiscent of that found by \citetalias{kewley06b} for the SDSS sources, mostly radio-quiet AGN."
 However. we fux a significant nuniber of LEG located above the line marking the transition between LINERS and Sevferts.," However, we find a significant number of LEG located above the line marking the transition between LINERs and Seyferts."
 The location of LEG shows an upward scatter with respect to the “finger” of highest LINERS deusity by ~0.2 dex iu the ΙΟ rratio., The location of LEG shows an upward scatter with respect to the `finger' of highest LINERs density by $\sim$ 0.2 dex in the [O ratio.
 As already noted there is a substantial nismatcl in 1niuositv between the 3CR aud 1e SDSS sources., As already noted there is a substantial mismatch in luminosity between the 3CR and the SDSS sources.
 Our data are not sufficient to conclude whether this is duc to a genuine difference between the (ostlv) racdio-quiet AGN of the SDSS aud the radio-loud AGN of the 3CR sample. (due e.g. to a contribution of jets emission to the line excitation) or simply to a Iuniünositv difference.," Our data are not sufficient to conclude whether this is due to a genuine difference between the (mostly) radio-quiet AGN of the SDSS and the radio-loud AGN of the 3CR sample, (due e.g. to a contribution of jets emission to the line excitation) or simply to a luminosity difference."
 Sinularh. we noted that the 3CR ποος are concentrated along the edges of the SDSS density distribution.," Similarly, we noted that the 3CR sources are concentrated along the edges of the SDSS density distribution."
" The first possibility. to account for the ocation of SDSS aud 3CR sources is again a general difference in the spectroscopic behavior between radio-quiet and radio-lou AGN,", The first possibility to account for the location of SDSS and 3CR sources is again a general difference in the spectroscopic behavior between radio-quiet and radio-loud AGN.
 Alternatively. he locatiou of he SDSS sources could be due to the contamination of star forming regious.," Alternatively, the location of the SDSS sources could be due to the contamination of star forming regions."
 In this scenario. this results in a arge spread in the line ratios. cepeuding ou the relative contribution of the line emission produced by the active micleus aud by star formation.," In this scenario, this results in a large spread in the line ratios, depending on the relative contribution of the line emission produced by the active nucleus and by star formation."
 The sources would be theu distributed along a mixing region (possibly the Προς. seen in Fies., The sources would be then distributed along a mixing region (possibly the `fingers' seen in Figs.
 1 and 2)). ranging from a ‘pure’ ACN to a star fornüng spectrum.," \ref{dd} and \ref{sdss4045}) ), ranging from a `pure' AGN to a star forming spectrum."
 Iu the case of 3C'R sources. the hieher line huninositv is likely to be imdicatiou of a dominant AGN contribution. also considering the ecucral," In the case of 3CR sources, the higher line luminosity is likely to be indication of a dominant AGN contribution, also considering the general"
mass or luminosity.,mass or luminosity.
" Therefore, it is interesting to check whether two slopes in SF are characteristic for the hole range of, say, black hole masses or have dependence on mass."," Therefore, it is interesting to check whether two slopes in SF are characteristic for the hole range of, say, black hole masses or have dependence on mass."
" Using black hole mass estimations from Shenetal. we select three subsamples of quasars with black hole (2008),,masses in ranges <5x105M (1953 sources, low mass range), 5x10°+1.510? (2108 sources, mean mass range), >1.5x10?Μο (1276M. sources, high mass and build SF for each range of masses."," Using black hole mass estimations from \cite{2008ApJ...680..169S}, we select three subsamples of quasars with black hole masses in ranges $\le5\times10^8\,M_\odot$ (1953 sources, low mass range), $5\times10^8\div
1.5\times10^9M_\odot$ (2108 sources, mean mass range), $\ge1.5\times10^9M_\odot$ (1276 sources, high mass range) and build SF for each range of masses."
" All three SFs range)show existence of two slopes (see Figure 8,, top panel)"," All three SFs show existence of two slopes (see Figure \ref{fig:sf_mbh}, top panel)."
 We fit SFs by broken power low model (Equation 3))., We fit SFs by broken power low model (Equation \ref{eq:bpl}) ).
" The best fit parameters and formal estimation of errors by ""jackknife' resampling are given in the Table 1..", The best fit parameters and formal estimation of errors by 'jackknife' resampling are given in the Table \ref{tab:fit_par}.
" In order to have better understanding of the uncertainties in the parameters we also plot regions showing minimal and maximal values of the fits for all three SFs Figure 8,, bottom panel)."," In order to have better understanding of the uncertainties in the parameters we also plot regions showing minimal and maximal values of the fits for all three SFs (see Figure \ref{fig:sf_mbh}, bottom panel)."
" As it is seen from the Figure(see 8,, the SFs slopes tend to increase with the growth of the mean subsample mass but still have common values."," As it is seen from the Figure \ref{fig:sf_mbh}, the SFs slopes tend to increase with the growth of the mean subsample mass but still have common values."
 More subtle study is needed in order to show that the trend in slopes as a function of black hole mass is real., More subtle study is needed in order to show that the trend in slopes as a function of black hole mass is real.
 Here we only conclude that the presence of two slopes in the ensemble SF is independent on the considered black hole mass range., Here we only conclude that the presence of two slopes in the ensemble SF is independent on the considered black hole mass range.
" We consider three processes frequently invoked to explain the variability of quasars: instabilities in the accretion flow around a supermassive black hole, supernova explosions related to the starburst phenomenon, and gravitational microlensing by compact bodies in the host galaxy (Kawaguchietal.1998;Hawkin"," We consider three processes frequently invoked to explain the variability of quasars: instabilities in the accretion flow around a supermassive black hole, supernova explosions related to the starburst phenomenon, and gravitational microlensing by compact bodies in the host galaxy \citep{1998ApJ...504..671K,2002MNRAS.329...76H}."
s SFs of light curves predicted by each model are 2002)..characterized by different slopes., SFs of light curves predicted by each model are characterized by different slopes.
" The expected SF slopes are 0.25+ 0.03, 0.44+0.03, 0.83+0.08 respectively for variability generated by miscrolensing, disk instability, and starburst mechanisms 2002).."," The expected SF slopes are $0.25\pm0.03$ , $0.44\pm0.03$, $0.83\pm0.08$ respectively for variability generated by miscrolensing, disk instability, and starburst mechanisms \citep{2002MNRAS.329...76H}."
 The SF slopes found in Section 4 are 0.79 (Hawkinsand 0.33 correspondingly for time lags below and above 42 days., The SF slopes found in Section \ref{sec:sf_ensemble} are 0.79 and 0.33 correspondingly for time lags below and above $42$ days.
" Taken at face value, the first slope is consistent with the starburst model, and the second slope falls half way between the predictions of the disk instability model and the microlensing model."," Taken at face value, the first slope is consistent with the starburst model, and the second slope falls half way between the predictions of the disk instability model and the microlensing model."
" Note that SF slopes in the range 0.30—0.35 have been obtained in other works, but the authors still attributed the variability to disk instabilities."," Note that SF slopes in the range 0.30–0.35 have been obtained in other works, but the authors still attributed the variability to disk instabilities."
 Baueretal.(2009) and Meusinger(2011) present an overview of recent results.," \cite{2009ApJ...696.1241B} and \cite{2011A&A...525A..37M}
 present an overview of recent results."
 Asymmetries in the SF provide an additional test of model predictions., Asymmetries in the SF provide an additional test of model predictions.
" The SF that only includes those pairs of measurements for which the flux increases with time, may be different from S_, the SF characterizing a decreaseS,, in flux."," The SF that only includes those pairs of measurements for which the flux increases with time, $S_+$, may be different from $S_-$, the SF characterizing a decrease in flux."
" Kawaguchietal. show that in starburst models S,>S_.", \cite{1998ApJ...504..671K} show that in starburst models $S_+>S_-$.
" Disk (1998)instabilities produce S_>S, and microlensing is symmetric, ie. when averaged over sufficiently long time intervals (Hawkins 2002).. "," Disk instabilities produce $S_->S_+$ and microlensing is symmetric, i.e. when averaged over sufficiently long time intervals \citep{2002MNRAS.329...76H}. ."
"The difference between S. and S, depends on the parameters of the model and therefore —ο is still possible in both the starburst and the disk instabilityS, model (Kawaguchietal."," The difference between $S_-$ and $S_+$ depends on the parameters of the model and therefore $S_-
\simeq S_+$ is still possible in both the starburst and the disk instability model \citep{1998ApJ...504..671K}."
" Noise-corrected S, and S_ functions are 1998)..shown in Figure 9 with blue and green lines respectively.", Noise-corrected $S_+$ and $S_-$ functions are shown in Figure \ref{fig:sf_plus_minus} with blue and green lines respectively.
" For long time lags in the range 300-1600 days, where the SF slope is —0.33, we have S_>Si, as predicted by the disk instability model."," For long time lags in the range 300–1600 days, where the SF slope is $\sim0.33$, we have $S_- > S_+$, as predicted by the disk instability model."
" The significance of the difference is not very high, but allows us to exclude the microlensing model."," The significance of the difference is not very high, but allows us to exclude the microlensing model."
" The difference disappears for time-scales below 100 days, where the slope is consistent with the starburst model."," The difference disappears for time-scales below 100 days, where the slope is consistent with the starburst model."
 This may suggest that a typical supernova rate in a starbursts in the host galaxies ofquasars is as high as ~100 γι! and effectively washes out any detectable asymmetry (see Fig., This may suggest that a typical supernova rate in a starbursts in the host galaxies ofquasars is as high as $\sim100$ $^{-1}$ and effectively washes out any detectable asymmetry (see Fig.
 6 in Kawaguchi (1998)))., 6 in \cite{1998ApJ...504..671K}) ).
" However, therate should be treated with"," However, therate should be treated with"
 (Schreieretal.L979.. (Schreicr Durusotal.1983.. Clarkeetal.1986)) Jonesetal.1996)). Blandford1991.. Autouncci1993)). Baileyctal.1986.. Packhametal.1996)). USTWE/PC-1tnagingpolarimetryoftheiuucr (Schreieretal.1996)) Iusetal.1991.. Ferrarese Macchettoetal.1997.. Bowerctal.1998))," \cite{schreier:79}, \cite{feigelson:81}) \cite{schreier:81}, \cite{burns:83}, \cite{clarke:86}) \cite{jones:96}) \cite{blandford:91}, \cite{antonucci:93}) \cite{bailey:86}, \cite{packham:96}) \cite{schreier:96}) \cite{harms:94}, \cite{ferrarese:96}, \cite{macchetto:97}, \cite{bower:98})"
origins of long-period ancl short-period comets are correct. MGU disk models make a clear prediction: long-period comets should be more isotopically homogeneous than short-period comets.,"origins of long-period and short-period comets are correct, MGU disk models make a clear prediction: long-period comets should be more isotopically homogeneous than short-period comets."
 In addition. both twpes of comet should contain relractory particles similar to CAIs. such as the Coki particle in Wild 2.," In addition, both types of comet should contain refractory particles similar to CAIs, such as the Coki particle in Wild 2."
 I thank the referee for a number of perceplive suggestions and Sandy Weiser for cluster and workstation management., I thank the referee for a number of perceptive suggestions and Sandy Keiser for cluster and workstation management.
 This research was supported in part by the NASA Planetary Geology and Geophysics Program (NNNOTADAGG) and. by the NASA Origins of Solar Systems Program (NNNOQAFG62C). and is contributed in part to the NASA Astrobiology Institute (NNAOO0DASLIA).," This research was supported in part by the NASA Planetary Geology and Geophysics Program (NNX07AP46G) and by the NASA Origins of Solar Systems Program (NNX09AF62G), and is contributed in part to the NASA Astrobiology Institute (NNA09DA81A)."
 Calenlations were perlormed on the Carnegie Alpha Cluster. which was supported in part by the NSF MBI Program (AST-9976615).," Calculations were performed on the Carnegie Alpha Cluster, which was supported in part by the NSF MRI Program (AST-9976645)."
"and we used and If Lm>Lo, then Lm is the comoving Jeans length and the Jeans mass is given by the solution of If Lm«Lo, then the average is done inside a single cell, using equation (7)), and the Jeans mass is given by We study primordial magnetic fields in the form of randomly oriented cells considering two possible scenarios for the seed field.","and we used and If $L_m>L_0$, then $L_m$ is the comoving Jeans length and the Jeans mass is given by the solution of If $L_m<L_0$, then the average is done inside a single cell, using equation \ref{eq:Brms2}) ), and the Jeans mass is given by We study primordial magnetic fields in the form of randomly oriented cells considering two possible scenarios for the seed field."
 The first scenario we discuss is one where each cell contains a dipole field whose flux is conserved., The first scenario we discuss is one where each cell contains a dipole field whose flux is conserved.
 In this case we have p—3/2 2010b) in equations (12)) and (13)).," In this case we have $p=3/2$ \citep{Hindmarsh1998,Souza2008,Souza2010} in equations \ref{eq:Lm}) ) and \ref{eq:Jeansfinal}) )."
" We also consider the geometry studied by Ahonen&Enqvist(1998) and Enqvist&Olesen(1993),, who found cells with large ring-like fields, but with planes of inclination randomly oriented."," We also consider the geometry studied by \citet{Ahonen1997} and \citet{Enqvist1993}, who found cells with large ring-like fields, but with planes of inclination randomly oriented."
" Thus, an average over large volumes corresponds to a random walk of all possible inclinations."," Thus, an average over large volumes corresponds to a random walk of all possible inclinations."
" This is equivalent a random walk on a 2D surface of a sphere, which implies p— 1."," This is equivalent a random walk on a 2D surface of a sphere, which implies $p=1$ ."
" In order to calculate the Jeans and filtering masses from equations (4)) and (13)), it is necessary to have an expression for the evolution of the temperature of the gas with redshift."," In order to calculate the Jeans and filtering masses from equations \ref{eq:filtering}) ) and \ref{eq:Jeansfinal}) ), it is necessary to have an expression for the evolution of the temperature of the gas with redshift."
" We use the analytic fit of the temperature as a function of redshift that Kravtsovetal.(2004) obtained for the results of Gnedin(2000),, where z>zs is the epoch before the first HII regions form, z,€zzs is the epoch of the overlap of multiple HIT regions and z<z, is the epoch of complete reionization."," We use the analytic fit of the temperature as a function of redshift that \citet{Kravtsov2004} obtained for the results of \citet{Gnedin2000a}, where $z > z_s$ is the epoch before the first HII regions form, $z_r \leq z \leq z_s$ is the epoch of the overlap of multiple HII regions and $z < z_r$ is the epoch of complete reionization."
" Throughout this paper we use a=6, z,=11 and Zp=8, unless otherwise mentioned."," Throughout this paper we use $\alpha=6$, $z_s=11$ and $z_r=8$, unless otherwise mentioned."
" We use equations (13)), (14)) and (15)) in (4)) to calculate the effect of RPMF on the filtering mass."," We use equations \ref{eq:Jeansfinal}) ), \ref{eq:Jeanshom}) ) and \ref{eq:temp}) ) in \ref{eq:filtering}) ) to calculate the effect of RPMF on the filtering mass."
" The results obtained by assuming different values for Lo and Bo are shown in figures 1 and 2,, for dipole-like fields, and in figures 5 and 6 for ring-like fields (without taking into account the effects of amplification, i.e. setting f7(z)zz 1)."," The results obtained by assuming different values for $L_0$ and $B_0$ are shown in figures \ref{fig:filter7} and \ref{fig:filter8}, for dipole-like fields, and in figures \ref{fig:filter7ring} and \ref{fig:filter8ring} for ring-like fields (without taking into account the effects of amplification, i.e. setting $f^2_{T}(z)\approx 1$ )."
 The model proposed by deSouza&Opher(2008) leads to dipole-like field with a comoving Boz0.14G and Lo©1pc., The model proposed by \citet{Souza2008} leads to dipole-like field with a comoving $B_0\approx 0.1\muG$ and $L_0\approx 1\text{ pc}$.
" This curve deviates only slightly from the case of no magnetic field, in figure 1.."," This curve deviates only slightly from the case of no magnetic field, in figure \ref{fig:filter7}."
" We found that most models where magnetic fields are generated during a quark-hadron phase transition — which would have dipole-like fields with Bo©2x1077G and Loz1A.U. (Hogan1983),, or Boz10716G and Loz1pc (Cheng&Olinto1994) — or during an electroweak phase transition — ring-like fields with Bo~107"" to 107?G and Lo~10A.U. (Baym1996) — have negligible effects on the filtering mass."," We found that most models where magnetic fields are generated during a quark-hadron phase transition – which would have dipole-like fields with $B_0\approx 2\times 10^{-17}\text{ G}$ and $L_0\approx 1\text{ A.U.}$ \citep{Hogan1983}, or $B_0\approx 10^{-16}\text{ G}$ and $L_0\approx 1\text{ pc}$ \citep{ChengOlinto} – or during an electroweak phase transition – ring-like fields with $B_0\sim 10^{-7}$ to $10^{-9}\text{ G}$ and $L_0\sim10\text{ A.U.}$ \citep{Baym1996} – have negligible effects on the filtering mass."
 Observations of the cosmic microwave background radiation (CMB) lead to an upper limit on the homogeneous primordial magnetic field Boys=2.98nG (comoving) (Yamazakietal.2010) with Lo~1Mpc.," Observations of the cosmic microwave background radiation (CMB) lead to an upper limit on the homogeneous primordial magnetic field $B_{CMB}=2.98\nG$ (comoving) \citep{Yamazaki2010} with $L_0\sim 1\,\text{Mpc}$."
" This limit corresponds to the brown curve plotted in figures 2,, 4,, 6 and 8.."," This limit corresponds to the brown curve plotted in figures \ref{fig:filter8}, \ref{fig:filter8turb}, \ref{fig:filter8ring} and \ref{fig:filter8ringturb}."
" There is, thus, a family of possible models to explain the origin of cosmic magnetic fields in the early Universe that can create a difference in the filtering mass between 104—109?Ma and is in agreement with the CMB constraints."," There is, thus, a family of possible models to explain the origin of cosmic magnetic fields in the early Universe that can create a difference in the filtering mass between $10^4-10^{9.5} M_{\odot}$ and is in agreement with the CMB constraints."
" 'The increase of the filtering mass due to the presence of magnetic fields is bigger before the reionization era, since the temperature, then, contributes less to the total pressure."," The increase of the filtering mass due to the presence of magnetic fields is bigger before the reionization era, since the temperature, then, contributes less to the total pressure."
 We also considered that the seed field could have been amplified by effects of intergalactic turbulence (as discussed in section 3.3))., We also considered that the seed field could have been amplified by effects of intergalactic turbulence (as discussed in section \ref{sec:turb}) ).
 The evolution of the filtering mass considering this effect is shown in figures 3 and 4 for dipole-fields and 7 and 8 for ring-like fields., The evolution of the filtering mass considering this effect is shown in figures \ref{fig:filter7turb} and \ref{fig:filter8turb} for dipole-likefields and \ref{fig:filter7ringturb} and \ref{fig:filter8ringturb} for ring-like fields.
" Comparing these figures with the previous ones, we note that the amplification leads to an increase in the filtering mass only at small redshifts."," Comparing these figures with the previous ones, we note that the amplification leads to an increase in the filtering mass only at small redshifts."
hayls)=1.92|2.727.,$k_{\rm av}(z) = 1.9z+2.7z^2$.
 This expression can be used in lieu of the above or instead. could. be applied. to non-classified galaxies which require Av -corrections., This expression can be used in lieu of the above or instead could be applied to non-classified galaxies which require $K$ -corrections.
 The most. substantial source of errors in our ÁA-corrections will arise. from the uncertainty in the 2dE system response., The most substantial source of errors in our $K$ -corrections will arise from the uncertainty in the 2dF system response.
 Taking this into account. we estimate the uncertainty in our A-corrections to be of the order of 20%., Taking this into account we estimate the uncertainty in our $K$ -corrections to be of the order of $20\%$.
 Norberg ((2001) have also calculated A -corrections for the 24ος5. independently of the observed galaxy spectra.," Norberg (2001) have also calculated $K$ -corrections for the 2dFGRS, independently of the observed galaxy spectra."
 Thes find very similar /v-corrections. which are consistent with our own to within the stated uncertainty.," They find very similar $K$ -corrections, which are consistent with our own to within the stated uncertainty."
" The total b, magnitudes used in this analysis have been derived from plate scans (Alackelox et al.", The total $\bj$ magnitudes used in this analysis have been derived from plate scans (Maddox et al.
 1990) and so they will be more susceptible to certain systematic ellects than those derived. from CCD data., 1990) and so they will be more susceptible to certain systematic effects than those derived from CCD data.
" In. order to improve the accuracy of our magnitudes several re-calibrations of the 5, magnitudes have been made using CCD photometry from overlapping fields (see Norberg ct al.."," In order to improve the accuracy of our magnitudes several re-calibrations of the $\bj$ magnitudes have been made using CCD photometry from overlapping fields (see Norberg et al.,"
 2001)., 2001).
 Unfortunately some small olfsets still remain., Unfortunately some small offsets still remain.
" One such ellect which may. be important to the analysis of Lbs per spectral twpe is an apparent shift in 5, magnitude with surface brightness (Cross 2001: Blanton et al.."," One such effect which may be important to the analysis of LFs per spectral type is an apparent shift in $\bj$ magnitude with surface brightness (Cross 2001; Blanton et al.,"
 private communication)., private communication).
 Fig., Fig.
" S shows the ollsets. between the 2PdbGRS b, magnitudes and those derived. from a CCD photometric survey (Cross 2001: Cross et al."," \ref{surf1} shows the offsets between the 2dFGRS $\bj$ magnitudes and those derived from a CCD photometric survey (Cross 2001; Cross et al.,"
 in preparation) versus surface brightness. a trend is clear in that high surface brightness objects tend to be systematically fainter in ο by.," in preparation) versus surface brightness, a trend is clear in that high surface brightness objects tend to be systematically fainter in 2dFGRS $\bj$."
 These olfsets. are most. likely due to a saturation effect in the (plate derived) 20EGIUS magnitudes., These offsets are most likely due to a saturation effect in the (plate derived) 2dFGRS magnitudes.
 These small olfsets will have a negligible elfect on the calculation of the overall LF but may be significant when one clivicles the 2drORS sample in a way which is related to surface brightness., These small offsets will have a negligible effect on the calculation of the overall LF but may be significant when one divides the 2dFGRS sample in a way which is related to surface brightness.
 Fie., Fig.
 9 shows how these shifts depend upon our spectral classification. it can be seen that the relationship is surprisingly weak.," \ref{surf2} shows how these shifts depend upon our spectral classification, it can be seen that the relationship is surprisingly weak."
 The mean shifts we calculate per type are. however significant compared. to. their estimated uncertaintics (Table. 22).," The mean shifts we calculate per type are however significant compared to their estimated uncertainties (Table. \ref{tabsb}) ),"
 particularly for the most [ate-tvpe galaxies., particularly for the most late-type galaxies.
 We correct the galaxy magnitudes for each spectral tvpe in order to account for these shifts in the analysis that follows., We correct the galaxy magnitudes for each spectral type in order to account for these shifts in the analysis that follows.
" We assume ao flat) homogeneous Universe with a uniform Llubble How ()=ffy/1l00 kkniss! ty, a cosmological constant ον=O7 ancl mass density"," We assume a flat homogeneous Universe with a uniform Hubble flow $h=H_0/100$ $^{-1}$ $^{-1}$ ), a cosmological constant $\Omega_\Lambda=0.7$ and mass density"
((2.2)).,\ref{eq:rho}) ).
" This density is integrated downward [rom some upper heieht vielding (he null column density. N,,(2). plotted in reffig:nd2-23.."," This density is integrated downward from some upper height yielding the null column density, $N_n(z)$, plotted in \\ref{fig:nd2-23}. ."
 The exponential factor. e.7. in each of the spectral integrals means that the null density al a given height 2 depends principally on the spectrum over wave numbers hoclfz.," The exponential factor, $e^{-2kz}$, in each of the spectral integrals means that the null density at a given height $z$ depends principally on the spectrum over wave numbers $k<1/z$."
 Thus heights above 2=1 Mm depend mostly on the spectrum below (to the left ol) &=Lrad/Mam.," Thus heights above $z=1$ Mm depend mostly on the spectrum below (to the left of) $k=1\,{\rm rad/Mm}$."
" As discussed above. ancl evident in rellie:spec,,/f. allM DIspectraare feirlysimilaroverthisrange."," As discussed above, and evident in \\ref{fig:spec_mtf}, all MDI spectra are fairly similar over this range."
 IHisthereforenaturalthatall[ivesolideuri 23followsimilarpathsaboveoneMm., It is therefore natural that all five solid curves in \\ref{fig:nd2-23} follow similar paths above one Mm.
 Noise contributes significant power to the highest wavenumbers. creating a population of artificial null points in an extrapolation.," Noise contributes significant power to the highest wavenumbers, creating a population of artificial null points in an extrapolation."
 The nullcolumn densities alter subtraction of this noise. the dotted curves in," The nullcolumn densities after subtraction of this noise, the dotted curves in"
cosmological simulations and. for the first time. we report on an adequate comparison between our expectations and observations.,"cosmological simulations and, for the first time, we report on an adequate comparison between our expectations and observations."
 The simulation was done using the cosmological simulation code GADGET-2 (2). with a treatment for magnetic fields., The simulation was done using the cosmological simulation code GADGET-2 \citep{2005MNRAS.364.1105S} with a treatment for magnetic fields.
 It eutures an entropy conserving formulation of Smooth Particle Hydrodynamics (SPH) (2)... which is supplemented with the ormulation of ideal MHD presented in ?)..," It features an entropy conserving formulation of Smooth Particle Hydrodynamics (SPH) \citep{2002MNRAS.333..649S}, which is supplemented with the formulation of ideal MHD presented in \citet{2008arXiv0807.3553D}."
 The implementation ‘ollows the induction equation and computes the back reaction of he magnetic field using a symmetric formulation of the Lorentz orce., The implementation follows the induction equation and computes the back reaction of the magnetic field using a symmetric formulation of the Lorentz force.
 We used a divergence cleaning scheme presented in 2).. which reduces numerical noise in shocks by subtracting the magnetic force which is proportional to the divergence of the field.," We used a divergence cleaning scheme presented in \citet{2001ApJ...561...82B}, which reduces numerical noise in shocks by subtracting the magnetic force which is proportional to the divergence of the field."
 It also helps to suppress the clumping instability particle based THD codes encounter in regions with small plasma «2 (.e. where magnetic pressure considerably exceeds thermal pressure)., It also helps to suppress the clumping instability particle based MHD codes encounter in regions with small plasma $\beta$ (i.e. where magnetic pressure considerably exceeds thermal pressure).
 In non radiative simulations like ours. regions with small jpasma 7 are rare.," In non radiative simulations like ours, regions with small plasma $\beta$ are rare."
 Only the strong shocks in cores of galaxy clusters during major mergers produce enough compression to amplify the tield to become dynamically dominant., Only the strong shocks in cores of galaxy clusters during major mergers produce enough compression to amplify the field to become dynamically dominant.
 These mergers are relatively brief events and are handled more accurately with our new numerical treatment (see.2.fordetails)., These mergers are relatively brief events and are handled more accurately with our new numerical treatment \citep[see ][ for details]{2008arXiv0807.3553D}.
 2?) have shown that non radiative simulations overpredict the gas density in cores of galaxy clusters., \citet{2006MNRAS.367.1641B} have shown that non radiative simulations overpredict the gas density in cores of galaxy clusters.
 This affects our simulation as well and can be seen in density and magnetic field profiles as well as in X-ray luminosities., This affects our simulation as well and can be seen in density and magnetic field profiles as well as in X-ray luminosities.
 As cosmological MHD SPH simulations lack physical dissipation. radiative SPH MHD simulations are not feasible at the moment.," As cosmological MHD SPH simulations lack physical dissipation, radiative SPH MHD simulations are not feasible at the moment."
 On the other hand secondary models have difficulties reproducing the outer parts of radio haloes correctly., On the other hand secondary models have difficulties reproducing the outer parts of radio haloes correctly.
 Therefore our main focus lies on these regions. where the simulations are not affected by the overpredicted gas density.," Therefore our main focus lies on these regions, where the simulations are not affected by the overpredicted gas density."
 We used a constrained realisation of the local universe (see ?) and referencestherein)., We used a constrained realisation of the local universe (see \citet{2005JCAP...01..009D} and referencestherein).
 The initial conditions are similar to those used in 2). to study the formation of the local galaxy population., The initial conditions are similar to those used in \citet{2002MNRAS.333..739M} to study the formation of the local galaxy population.
 They were obtained based on the IRAS 1.2-]y galaxy survey., They were obtained based on the IRAS 1.2-Jy galaxy survey.
 Its density field was smoothed on a scale of 7Mpc. evolved back in time to 2=50 using the Zeldovich approximation and assumed to be Gaussian (2)..," Its density field was smoothed on a scale of $7\, \mathrm{Mpc}$, evolved back in time to $z=50$ using the Zeldovich approximation and assumed to be Gaussian \citep{Hoffman1991}."
 The IRAS observations constrain a volume of zz115Mpc centered on he Milky Way.," The IRAS observations constrain a volume of $\approx 115 \, \mathrm{Mpc}$ centered on the Milky Way."
 It was sampled with dark matter particles and embedded in a periodic box of z343Alpe comoving.," It was sampled with dark matter particles and embedded in a periodic box of $\approx 343 \, \mathrm{Mpc}$ comoving."
 Outside of the inner region. the box is filled with dark matter particles with 1/6t 1ofthe resolution. to cover for long range gravitational tidal forces arising from the low-frequency constrains.," Outside of the inner region, the box is filled with dark matter particles with $1/6$ th of the resolution, to cover for long range gravitational tidal forces arising from the low-frequency constrains."
 In the evolved density Ποd. many locally observed galaxy clusters can be identified by position and mass.," In the evolved density field, many locally observed galaxy clusters can be identified by position and mass."
 Especially the Coma cluster (see2) shows remarkable similarities in morphology., Especially the Coma cluster \citep[see][]{2010MNRAS.401...47D} shows remarkable similarities in morphology.
 A fly-through of the simulation can be downloaded from the MPA The initial conditions were extended to include gas by splitting dark matter particles in the high resolution region into gas and dark matter particles of masses 0.69.LO?AL. and 44.10°M. respectively.," A fly-through of the simulation can be downloaded from the MPA The initial conditions were extended to include gas by splitting dark matter particles in the high resolution region into gas and dark matter particles of masses $0.69 \times 10^9\; {\rm M}_\odot$ and $4.4
\times 10^9\; {\rm M}_\odot$ respectively."
 Therefore the biggest clusters are resolved by about a million particles., Therefore the biggest clusters are resolved by about a million particles.
 The gravitational softening length was set to LOkpe.," The gravitational softening length was set to $10\,\mathrm{kpc}$."
 This is comparable to the inter-particle separation found in the centre of the largest clusters., This is comparable to the inter-particle separation found in the centre of the largest clusters.
 The origins of magnetic fields in galaxy clusters are still under debate., The origins of magnetic fields in galaxy clusters are still under debate.
 It is assumed that some kind of early seed magnetic field is amplitied by structure formation through adiabatic compression. turbulence and shear flows to values observed today (1.10iC: in clusters).," It is assumed that some kind of early seed magnetic field is amplified by structure formation through adiabatic compression, turbulence and shear flows to values observed today $\approx 1 - 10
\,\mu\mathrm{G}$ in clusters)."
 Three main classes of models for the seed field exist: A first the seed fields can be created in shocks through the “Biermann battery 7 (222)...," Three main classes of models for the seed field exist: At first the seed fields can be created in shocks through the ”Biermann battery ” \citep{1997ApJ...480..481K,Ryu..1998,2001ApJ...562..233M}."
 A second class of models invokes primordia processes to predict seed fields that fill the entire volume of the universe., A second class of models invokes primordial processes to predict seed fields that fill the entire volume of the universe.
 The coherence length of these fields strongly depends on the details of the model (see.2.forareview)..," The coherence length of these fields strongly depends on the details of the model \citep[see
][ for a review]{Grasso..PhysRep.2000}."
 Finally the seec can be produced by AGN (22) or starbursting galaxies (2). at high redshift (22:4.— 6). whose outflows contaminate the proto-cluster Cosmological simulations using SPH (222). and grid based Adaptive Mesh Retinement (AMR) codes (22?) were able to show that observed Faraday rotations are compatible with a cosmological seed field of zz10.1!G.," Finally the seed can be produced by AGN \citep{1997ApJ...477..560E,2001ApJ...556..619F} or starbursting galaxies \citep{Volk&Atoyan..ApJ.2000} at high redshift $z \approx 4 - 6$ ), whose outflows contaminate the proto-cluster Cosmological simulations using SPH \citep{1999A&A...348..351D,2002A&A...387..383D,2005JCAP...01..009D}
 and grid based Adaptive Mesh Refinement (AMR) codes \citep{2005ApJ...631L..21B,2008A&A...482L..13D,2008ApJS..174....1L}
 were able to show that observed Faraday rotations are compatible with a cosmological seed field of $\approx 10^{-11}\,\mathrm{G}$."
 They also suggest that spatial distribution and structure of cluster magnetic fields are determined by the dynamics in the velocity field caused by structure formation (22).," They also suggest that spatial distribution and structure of cluster magnetic fields are determined by the dynamics in the velocity field caused by structure formation \citep{1999A&A...348..351D,2002A&A...387..383D}."
 For this work we follow ?) in terms of magnetic field origin., For this work we follow \citet{2008arXiv0808.0919D} in terms of magnetic field origin.
 They use a semianalytic mode for galactic winds (2). to seed magnetic fields in a constrained cosmological MHD SPH simulation., They use a semianalytic model for galactic winds \citep{2006MNRAS.370..319B} to seed magnetic fields in a constrained cosmological MHD SPH simulation.
 The continuous seeding process is approximated with an instantaneous seed at 2zd., The continuous seeding process is approximated with an instantaneous seed at $z\approx 4$.
 Astyey were able to show. the main properties of magnetic fields obained in clusters were not influenced by that approximation.," As they were able to show, the main properties of magnetic fields obtained in clusters were not influenced by that approximation."
 The wind model used assumes adiabatic expansion of a spherical gas bubble with homogeneous magnetic energy density around every galaxy below a certain mass threshold., The wind model used assumes adiabatic expansion of a spherical gas bubble with homogeneous magnetic energy density around every galaxy below a certain mass threshold.
 The magnetic bubble can be characterised by radius and tield strength., The magnetic bubble can be characterised by radius and field strength.
 The galaxy injects gas into the bubble carring frozen-in magnetic field from the dise into the bubble over the star-burst timescale., The galaxy injects gas into the bubble carring frozen-in magnetic field from the disc into the bubble over the star-burst timescale.
 Its final size is determined by the wind velocity. which is a function of the star formation rate and the properties of the ISM.," Its final size is determined by the wind velocity, which is a function of the star formation rate and the properties of the ISM."
 ?) give an evolution equation for the magnetic energy in the bubble depending on the star-burst timescale., \citet{2006MNRAS.370..319B} give an evolution equation for the magnetic energy in the bubble depending on the star-burst timescale.
 The energy is converted into a dipole moment and seeded once at a chosen redshift., The energy is converted into a dipole moment and seeded once at a chosen redshift.
 The magnetic field is then amplitied by structure formation to //€i level., The magnetic field is then amplified by structure formation to $\mu \mathrm{G}$ level.
 For details on the wind model refer to ??).. ," For details on the wind model refer to \citet{2006MNRAS.370..319B,2008arXiv0808.0919D}. ."
Figure |. shows full sky maps produced from the simulation. projecting the electron density. temperature," Figure \ref{vis} shows full sky maps produced from the simulation, projecting the electron density, temperature"
directions in kspace.,directions in ${\bf k}-$ space.
 This is done to eusure that the the field has a random phase., This is done to ensure that the the field has a random phase.
 For exiuuple. for the 4- conrponeut. we set up a complex field such that where G ids a function of a uniformly distributed independent random variables 44 or 0» that returus CGaussian-distributed raudoni values.," For example, for the $x$ -component, we set up a complex field such that where $G$ is a function of a uniformly distributed independent random variables $u_{1}$ or $u_{2}$ that returns Gaussian-distributed random values."
" The amplitude 2 1s given by where kh,=Ὁπίλ, aud AQ~Lh| kpe is a snoothing wavelength.", The amplitude $B$ is given by where $k_{o}=2\pi/\lambda_{o}$ and $\lambda_{o}\sim 43h^{-1}$ kpc is a smoothing wavelength.
 We then apply a divergence cleaning operator in k space as follows: where k is the unit vector m k-space., We then apply a divergence cleaning operator in ${\bf k}-$ space as follows: where $\hat{\bf k}$ is the unit vector in ${\bf k}$ -space.
 Note that this operator docs not change the shape of the power spectrum of the maguetic field. fluctuations., Note that this operator does not change the shape of the power spectrum of the magnetic field fluctuations.
 This is suffiicicut for our purposes as we assume the field is dynamically unimportant., This is sufficient for our purposes as we assume the field is dynamically unimportant.
 We then use threc-dimensional iuverse fast Fourier transform to convert the field in k-space to the real space., We then use three-dimensional inverse fast Fourier transform to convert the field in ${\bf k}$ -space to the real space.
 We solve the MIID equations. including field-aligued thermal conduction transport.," We solve the MHD equations, including field-aligned thermal conduction transport."
 These equations where where pry is the eas pressure and e is the gas internal energv per unit volume., These equations where where $p_{\rm th}$ is the gas pressure and $\epsilon$ is the gas internal energy per unit volume.
 We assuued the adiabatic index ~=4/3., We assumed the adiabatic index $\gamma = 5/3$.
 The anisotropic thermal conduction heat fiux F is even bx where ep is a undt vector pointing in the direction of the magnetic field and «# is the Spitzer-Dragiuskii ↸⊳∪∐≼⊔∐⊳⊓∪∐↸⊳∪↸∖≸−∏↸⊳↕↸∖∐↑∶↴∙⊾↕↖↽↸∖∐↴⋝∙↖↽∕⋅⊳∶⊓⊳∖↓∩∣⊺⊽↾−↸∖↥⋅∶↴∙⊾⋅ ⋅⋅ ⋅ ⋅↽−↣⋅≽ ↴∖↴↓↸⊳⋯↓↕↘⊽↓∙ ↕∐↸∖≺∣∏⋜↧↑↕∪∐⊔⊓∙⊂⊺↥⋅↸∖↻↥⋅↸∖↴∖↴↸∖∐↑↴∖↴↑∐↸∖↸⊳∪∪∐∐∶↴∙⊾↥⋅⋜↧↑↸∖↻↸∖↥⋅∏∐↕↑ ↖↽∪↕⋯⊔↸∖∙↖↖⊽↸∖∏↴∖↴↸∖↴," The anisotropic thermal conduction heat flux ${\bf F}$ is given by where $\hat{\bf e}_{B}$ is a unit vector pointing in the direction of the magnetic field and $\kappa$ is the Spitzer-Braginskii conduction coefficient given by $\kappa = 4.6\times 10^{-7}T^{5/2}$ erg $^{-1}$ $^{-1}$ $^{-1}$ In equation (21), ${\cal C}$ represents the cooling rate per unit volume."
∖↴↑⋜∐∐↧⋜∐⋅≼⇂↑⋜∏⋝∏↕⋜↧↑↸∖≼⇂↻∏↴⋝∐↸⊳↕⋅↖↽⋜↧↖↽⋜↧∐⋜∏⋝↕↸∖ ↸⊳∪∪∐∐∶↴∙⊾↸⊳↿∐⋅↖↽↸∖↴∖↴⋖∎∙↗⋟↕⋟∪↥⋅⋯↸∖↑⋜↕∐↕↸⊳↕↑⋅↖↽∑∶∩∙∶≩∑∙∙ ⋀∖↕⋜↧∶↴∙⊾∐↸∖↑↕, We use standard tabulated publicly available cooling curves \citep{sutherland93} for metallicity $Z = 0.3Z_{\odot}$.
↸⊳∱∎∐∖∐↸∖↖↽∪↕∏↕∪∐↖↖↽⋜↧↴∖↴↴∖↴∪↕↖↽↸∖≺⇂↴⋝∙↖⇁⋯↸∖⋜⋯↴∖↴ ∪↕⋟⋜↧≼∐↥⋅↸∖↸⊳↑↕∪∐⋜↧∐⋅↖⇁∏∐↴∖↴↻↕↑↴∖↴↑⋜↧∶↴∙⊾∶↴∙⊾↸∖↥⋅↸∖≼↧⋯↸∖↴∖↴∐⋜↧↕∶↴⋁∪↥⋅↕↑∐⋯ (USA: Lee Deane 2009)., Magnetic field evolution was solved by means of a directionally unsplit staggered mesh algorithm (USM; Lee Deane 2009).
 The USAI module is based on a fnite-volune. high-order Godunuov scheme colmbined with coustrained transport method (CT).," The USM module is based on a finite-volume, high-order Godunov scheme combined with constrained transport method (CT)."
 This approach enarautees civergeuce-free magnetic fick distribution., This approach guarantees divergence-free magnetic field distribution.
 We implemented the anisotropic conduction unit following the approach of Sharma Waunet (2007)., We implemented the anisotropic conduction unit following the approach of Sharma Hammet (2007).
 More specifically. we applied monotonized ceutra (AIC) limiter to the coucuetive fluxes.," More specifically, we applied monotonized central (MC) limiter to the conductive fluxes."
 This methoc ensures that anisotropic conduction does not lear to negative temperatures in the presence of steep temperature eradicuts., This method ensures that anisotropic conduction does not lead to negative temperatures in the presence of steep temperature gradients.
 Tests of the method are presente ii acconipauviug paper (BRuszkowski et al., Tests of the method are presented in accompanying paper (Ruszkowski et al.
 2009. in preparation).," 2009, in preparation)."
 Tn order to emulate the effects of turbulence iu the ICAL we emploved a spectral forcing scheme that enables statistically stationary velocity fields (Eswaran Pope 1988).," In order to emulate the effects of turbulence in the ICM, we employed a spectral forcing scheme that enables statistically stationary velocity fields (Eswaran Pope 1988)."
 This scheme utilizes au Orustein-Ulleubeck random process. analogous to Browmiau motions iu a viscous inediun.," This scheme utilizes an Ornstein-Uhlenbeck random process, analogous to Brownian motions in a viscous medium."
 At cach timestep. accelerations are applied to the eas.," At each timestep, accelerations are applied to the gas."
 The acceleration field has vanishing mean aud cach acceleration mode has constant dispersion and is time-correlated., The acceleration field has vanishing mean and each acceleration mode has constant dispersion and is time-correlated.
 That is. at cach timestep the acceleration in a given direction is a Gaussian random variable with a fixed amplitude aud decays with time as XvalΈλι where f=exptt/Taceay).," That is, at each timestep the acceleration in a given direction is a Gaussian random variable with a fixed amplitude and decays with time as $\propto \sqrt{1-f^{2}}$, where $f=\exp (-t/\tau_{\rm decay})$."
 This essentially means that the forcing terii has a “memory of the previous state., This essentially means that the forcing term has a “memory” of the previous state.
 The phases are evolved in Fourier space and the divergence in the Sow is cleaned making the flow incompressible. V:e=0.," The phases are evolved in Fourier space and the divergence in the flow is cleaned making the flow incompressible, $\nabla \cdot v =0$."
 This is consistent with the Boussinesq approximation. aud justified since the eas motions are significantlv subsonic.," This is consistent with the Boussinesq approximation, and justified since the gas motions are significantly subsonic."
 Further details aud the nunerical tests of this method cau be found im Fisher et al. (, Further details and the numerical tests of this method can be found in Fisher et al. (
2008).,2008).
 The kev quautity of interest is the velocity dispersion of the resultant velocity field σ and its scaling with the injected energy aud the injection spatial scale., The key quantity of interest is the velocity dispersion of the resultant velocity field $\sigma$ and its scaling with the injected energy and the injection spatial scale.
 This can be understood iu terms of simple dimensional analysis in an equilibrium situation when the turbulent diving is balanced by viscous dissipation in the fluid where ἁγμοας Is the number of Fourier driving nodes. € is the energv injection per wit mass per mode aud poc Leis the gas viscosity for aa eiven larecst eddy," This can be understood in terms of simple dimensional analysis in an equilibrium situation when the turbulent driving is balanced by viscous dissipation in the fluid where $N_{\rm mode}$ is the number of Fourier driving modes, $\epsilon^{*}$ is the energy injection per unit mass per mode and $\nu\sim L\upsilon$ is the gas viscosity for a a given largest eddy"
That the best-fit temperatures in the two temperature model fits to the low and the high mass loss rate simulations are so similar. with only the relative normalisation of the two components varving. suggests the leniperatures are determined more by theZZOSAT PSPC's response than the true temperature distribution of the source.,"That the best-fit temperatures in the two temperature model fits to the low and the high mass loss rate simulations are so similar, with only the relative normalisation of the two components varying, suggests the temperatures are determined more by the PSPC's response than the true temperature distribution of the source."
 From Fig., From Fig.
 5. it is clear that for ZZ0.7keV the emission measure is well represented by a power law in 7. with a slope ofzLS.," \ref{fig:em} it is clear that for $T \ltsimm 0.7 \keV$ the emission measure is well represented by a power law in $T$, with a slope of $\gamma \approx -1.8$."
 Between 1=0.LkeV and 2=0.7keV the true emission measure falls by nearly two orders of magnitude. so a differential emission measure model with this slope should provide a good approximation to the spectrum.," Between $T = 0.1 \keV$ and $T = 0.7
\keV$ the true emission measure falls by nearly two orders of magnitude, so a differential emission measure model with this slope should provide a good approximation to the spectrum."
 The results of the model fits (Lable 6)) are therefor surprising. given the best fit values for >=0.25!," The results of the model fits (Table \ref{tab:mdm_fits}) ) are therefor surprising, given the best fit values for $\gamma \approx -0.25$!"
 The combination of material hotter than Z7—LkeV and the energy. dependent response of theROSALT PSPC must bias the fit significantly., The combination of material hotter than $T \sim 1 \keV$ and the energy dependent response of the PSPC must bias the fit significantly.
 The slope is clearly. wrong within the enerev range 0.1-2.4keV. ROSAT is sensitive to. let alone extended to the cut-olf temperatures claimed by the fit.," The slope is clearly wrong within the energy range $0.1$ $2.4 \keV$ ROSAT is sensitive to, let alone extended to the cut-off temperatures claimed by the fit."
 Are any of the effects above due to low numbers of photons?, Are any of the effects above due to low numbers of photons?
 Repeating the above analysis with assumed exposure times of 100008. giving simulated: spectra. with 3000 counts. eives results very similar to those in Section 3..," Repeating the above analysis with assumed exposure times of $10000 \s$, giving simulated spectra with $\gtsimm 3000$ counts, gives results very similar to those in Section \ref{sec:results}."
 The confidence regions quoted on the incliviclual fits are smaller. but the average best fit result is consistent. with those quoted above. and are statistically acceptable fits.," The confidence regions quoted on the individual fits are smaller, but the average best fit result is consistent with those quoted above, and are statistically acceptable fits."
 Lt therefor appears that the systematic deviations of the spectral fit results from the true values is due to fitting an intrinsicallv complex spectrum with a simplistic mocdel. and not due to poor photon statistics.," It therefor appears that the systematic deviations of the spectral fit results from the true values is due to fitting an intrinsically complex spectrum with a simplistic model, and not due to poor photon statistics."
 Can we infer any of the truc bubble properties from the spectral fits?, Can we infer any of the true bubble properties from the spectral fits?
 In particular. the true luminosity. densities. oressures and thermal energy are interesting quantities that we would like to know in addition to the temperature if observations are to be of any use in understanding he object.," In particular, the true luminosity, densities, pressures and thermal energy are interesting quantities that we would like to know in addition to the temperature if observations are to be of any use in understanding the object."
 The 0.1-2.4keV. fluxes predicted. from all the models are reasonably accurate (compare the real values from Table 2 with the results from he spectral fitting. Tables 4. - 6)).," The $0.1$ $2.4 \keV$ fluxes predicted from all the models are reasonably accurate (compare the real values from Table \ref{tab:bub_properties} with the results from the spectral fitting, Tables \ref{tab:rz_fits} - \ref{tab:mdm_fits}) )."
 The single temperature models get the intrinsic luminosity very wrong. due to the incorrect. absorption column. either overestimating it by a actor 5 (metallicity fixed. at solar) or underestimating ov an order of magnitude in the case of the fits with the metallicity. free. to fit.," The single temperature models get the intrinsic luminosity very wrong, due to the incorrect absorption column, either overestimating it by a factor $\sim5$ (metallicity fixed at solar) or underestimating by an order of magnitude in the case of the fits with the metallicity free to fit."
 On average the two temperature models overestimate the intrinsic flux by an order of magnitude. although. with a Large uncertainty due to the large variation in best-fit column.," On average the two temperature models overestimate the intrinsic flux by an order of magnitude, although with a large uncertainty due to the large variation in best-fit column."
 Similarly the cilferential emission measure models also have a large scatter in inferred intrinsic luminosity., Similarly the differential emission measure models also have a large scatter in inferred intrinsic luminosity.
 Assuming we know the volume of the emitting plasma V. the root mean square electron. density is ne= 1. where ΑΙ is the (volume) emission. measure emission measure obtained [rom the spectral fit.," Assuming we know the volume of the emitting plasma $V$, the root mean square electron density is $n_{\rm e} = \sqrt(EM/V)$ , where $EM$ is the (volume) emission measure emission measure obtained from the spectral fit."
 Phen the thermal pressure ancl energy. are. P?zz20KT. and. fp AÁTV.," Then the thermal pressure and energy are $P
\approx 2 n_{\rm e} k T$ and $E_{\rm TH} \approx 3 n_{\rm e} k T V$ ."
 Given the X-ray surface. brightness (Fig. 10)), Given the X-ray surface brightness (Fig. \ref{fig:ximage}) )
 we can estimate the volume., we can estimate the volume.
" The inferred properties are relatively insensitive to V as P and n», are proportional to V.(7 and Lynxyt27", The inferred properties are relatively insensitive to $V$ as $P$ and $n_{\rm e}$ are proportional to $V^{-1/2}$ and $E_{\rm TH} \propto V^{1/2}$.
 The bubble is clearly [limb-brishtened. (Fig. 10)).," The bubble is clearly limb-brightened (Fig. \ref{fig:ximage}) ),"
 the emission coming predominantly [from a thin shell near the shocked-wind/cold. shell interface., the emission coming predominantly from a thin shell near the shocked-wind/cold shell interface.
" Erom a radial profile of heROSAT surface brightness. and correcting for the point spread function. the inferred radial thickness of this shell is ~20"". the PSE-corrected. volume Vo;=OSI0""em."," From a radial profile of the surface brightness, and correcting for the point spread function, the inferred radial thickness of this shell is $\sim 20''$, the PSF-corrected volume $V_{\rm corr} = 9.8 \times 10^{56} \cm^{3}$."
" Dv wav of comparison. the total volume within the bubble inferred [rom the X-ray image is Viu=6.00510em""."," By way of comparison, the total volume within the bubble inferred from the X-ray image is $V_{\rm bub} = 6.0 \times 10^{57} \cm^{3}$."
 Note hat this is less than the volume quoted in Table 2. as we are not including the cold shell., Note that this is less than the volume quoted in Table \ref{tab:bub_properties} as we are not including the cold shell.
 The density. pressure and thermal energy inferred from. he single temperature models. using Yo; and Vin. are given in Table 7..," The density, pressure and thermal energy inferred from the single temperature models, using $V_{\rm corr}$ and $V_{\rm bub}$, are given in Table \ref{tab:inferred}."
 For the two temperature models. it would. be sensible o associate the cool component (with a high emission measure) of the spectral fit. to cool. dense material near he shocked-wind/cold shell interface. given the observed imb brightening.," For the two temperature models, it would be sensible to associate the cool component (with a high emission measure) of the spectral fit, to cool, dense material near the shocked-wind/cold shell interface, given the observed limb brightening."
 The hot component in the fit. with a total emission measure two to three orders of magnitude less than he cool component. could well represent emission from the rotrarefied gas in the bubble interior.," The hot component in the fit, with a total emission measure two to three orders of magnitude less than the cool component, could well represent emission from the hotrarefied gas in the bubble interior."
 Assigning volumes, Assigning volumes
Marri&Ferrara(1998) have discussed the feasibility of using gravitational lensing magnification due to intervening mass concentrations to observe high redshilt supernovae more easily. and potentially bring even twpe II supernovae within reach of NGST at 10.,"\citet{MarFer98} have discussed the feasibility of using gravitational lensing magnification due to intervening mass concentrations to observe high redshift supernovae more easily, and potentially bring even type II supernovae within reach of NGST at $z\ga 10$ ."
 Similar considerations also apply to our caleulations. and (his may help in obtaining detailed observations of PISNe.," Similar considerations also apply to our calculations, and this may help in obtaining detailed observations of PISNe."
 It is important to be able to identilv them reliably. distinguish them from other (vpes of supernova. and measure their redshift if we are to use them [for cosmological purposes.," It is important to be able to identify them reliably, distinguish them from other types of supernova, and measure their redshift if we are to use them for cosmological purposes."
 To be able to do this. we need a sample of SNe that are bright enough for spectroscopic observations.," To be able to do this, we need a sample of SNe that are bright enough for spectroscopic observations."
 Since none have been observed locally. we need the additional lensing magnification to investigate SNe in detail at high redshift.," Since none have been observed locally, we need the additional lensing magnification to investigate SNe in detail at high redshift."
 If ow mechanism for producing extvemely low metallicity. low mass stars (Pop 1I.5 in our lerminologv) via shock compression occurs in reality. then some of these stars should still be present in the halo of our galaxy.," If our mechanism for producing extremely low metallicity, low mass stars (Pop II.5 in our terminology) via shock compression occurs in reality, then some of these stars should still be present in the halo of our galaxy."
 Thev would of course be very okd (~14 GGyr). so only the lowest mass stars with M<0.8M. would remain (Girardi&Lallanzio 2002).," They would of course be very old $\sim 14$ Gyr), so only the lowest mass stars with $M\la 0.8\;\msun$ would remain \citep*{Giretal00,SieLivLat02}."
. We can use our model to predict how many such stars we would expect to find in a (vpical Milkv. Waa-sized halo anc. given a model [for the stellar halo density profile. (heir imber densitv in the solar neighborhood (we use a similar approach to that ol Hernandez&Ferrara 2001)).," We can use our model to predict how many such stars we would expect to find in a typical Milky Way-sized halo and, given a model for the stellar halo density profile, their number density in the solar neighborhood (we use a similar approach to that of \citealt{HerFer01}) )."
 Significant uncertainties in this ealeulation are the efficiency of producing Pop IL5 stars (as discussed in relsec:pop2.5)). the fraction of them that end up in the halo as opposed to the bulge. and iheir mass function.," Significant uncertainties in this calculation are the efficiency of producing Pop II.5 stars (as discussed in \\ref{sec:pop2.5}) ), the fraction of them that end up in the halo as opposed to the bulge, and their mass function."
 The star formation rate of Pop IL5 stars is given bv (u(2:11L5)=gel) in the redshift range when they form.," The star formation rate of Pop II.5 stars is given by $\psi_*(z; {\rm
II.5}) = \eta \psi_*(z; {\rm III})$ in the redshift range when they form."
 There is a slight difference to the wav οHI) is calculated for this problem. in that we are interested in the stars that end up in the Milkv Way (MW) halo.," There is a slight difference to the way $\psi_*(z; {\rm III})$ is calculated for this problem, in that we are interested in the stars that end up in the Milky Way (MW) halo."
 Pherefore. instead of the usual Press-Schechter mass function in the model. we use the Extended Press-Schechter mass funetion of halos at redshift z which will end up in a halo the mass of the MW at z=0 (Lacey&Cole1993).," Therefore, instead of the usual Press-Schechter mass function in the \citet{SanBroKam02} model, we use the Extended Press-Schechter mass function of halos at redshift $z$ which will end up in a halo the mass of the MW at $z=0$ \citep{LacCol93}."
". For a MW mass (total) of Mai=LO’M, we find that 10M. of Pop IHE stars are formed in MW progenitor halos. over the redshift range from z=30 to 2=15 in which we expect Pop 11.5 stars to form."," For a MW mass (total) of $M_{\rm MW}=10^{12}\;\msun$ we find that $10^{7.5}\;\msun$ of Pop III stars are formed in MW progenitor halos, over the redshift range from $z=30$ to $z=15$ in which we expect Pop II.5 stars to form."
 Interestingly. becausethe metal vield of PISNe is  0.5. this results in a MW metal enrichment of," Interestingly, becausethe metal yield of PISNe is $\sim 0.5$ , this results in a MW metal enrichment of"
one of the radii of curvature of the CME.,one of the radii of curvature of the CME.
 The underlving structure of a CME is believed to be a Πακ rope. which has two characteristic curvatures: a smaller one clue the curvature perpendicular to its axis (the radius when viewed as a cross-section) and the larger curvature along (he axis.," The underlying structure of a CME is believed to be a flux rope, which has two characteristic curvatures; a smaller one due the curvature perpendicular to its axis (the radius when viewed as a cross-section) and the larger curvature along the axis."
 —1 this Letter. we investigate if the shock relations above hold for a CAIE-driven IP shock.," In this Letter, we investigate if the shock relations above hold for a CME-driven IP shock."
 Specifically. we use direct observations of a CATE-driven shock observed in COR? and ]II1 instruments of the Sun Earth Connection Coronal ancl IHeliospheric Investigation suite (SECCII. Towardοἱal. 2008)) on (Ixsiserοἱal.2003).," Specifically, we use direct observations of a CME-driven shock observed in COR2 and HI1 instruments of the Sun Earth Connection Coronal and Heliospheric Investigation suite (SECCHI, \citealt{Howard:2008p4742}) ) on \citep{Kaiser:2008p1663}."
. In Section 2.. we present SECCILI observations of the CME and resulting shock and describe the analysis techiinq(que.," In Section \ref{s_obs}, we present SECCHI observations of the CME and resulting shock and describe the analysis technique."
 The results of our analvsis are presented in Section 3.., The results of our analysis are presented in Section \ref{s_res}.
 We discuss our results aud state our conclusions in Section 4 The CME analvsed first appeared in the COR (Thompsonetal.2003). coronagraph images [rom at UUT on 2008 April 5.," We discuss our results and state our conclusions in Section \ref{s_disc}
 The CME analysed first appeared in the COR1 \citep{Thompson:2003p1587} coronagraph images from at UT on 2008 April 5."
 It was most likely associated with a D-class flare from NOAA active region 10987. which was just behind the west limb as viewed [rom Earth.," It was most likely associated with a B-class flare from NOAA active region 10987, which was just behind the west limb as viewed from Earth."
" Figure 2. shows the CME as it propagates oul [rom the Sun into the different instruments field-ol-views [rom δα, to RR...", Figure \ref{f1} shows the CME as it propagates out from the Sun into the different instruments' field-of-views from $_{\odot}$ to $_{\odot}$.
 The CME was visible in both andB spacecralt in the inner and outer coronagraphs (COR and COR2: see movie MI). but was only visible in HI] (Evlesetal.2000). fromB (see movie AI2).," The CME was visible in both and spacecraft in the inner and outer coronagraphs (COR1 and COR2; see movie M1), but was only visible in HI1 \citep{Eyles:2009p3861} from (see movie M2)."
" The CME propagation direction was [ος to be 106"" west of the Sun-Earth line. the spacecralt were al a separation angle of 48° degrees (Irom each other)."," The CME propagation direction was found to be $\sim$ $^{\circ}$ west of the Sun-Earth line, the spacecraft were at a separation angle of $^{\circ}$ degrees (from each other)."
 The shock is visible as à curved brightness enhancement in both the COR? images in Figure 2. ancl also the Il]P image., The shock is visible as a curved brightness enhancement in both the COR2 images in Figure \ref{f1} and also the HI1 image.
 The accompanying movies M1 and M2 show the shock more clearly., The accompanying movies M1 and M2 show the shock more clearly.
 We have made (he assumption that the curved [ront is a shock. there are no radio orin-sifu data available to corroborate this.," We have made the assumption that the curved front is a shock, there are no radio or data available to corroborate this."
 However. due to the CME's velocity (1000 !) and the smoothness aud position of the feature ahead of the CME. it can be argued that (his it is a legitimate assumption (Bemporad&Mancuso2010:OntiverosVourlidas2009).," However, due to the CME's velocity $\sim$ $^{-1}$ ) and the smoothness and position of the feature ahead of the CME, it can be argued that this it is a legitimate assumption \citep{Bemporad:2010p9404, Ontiveros:2009p8787}."
. The observations were reduced using fom theSOLARSOFT library IIandsy 1998)..., The observations were reduced using from the library \citep{Freeland:1998p3546}.
 This corrects for a number of effects such as bias. Hlat-field ancl distortions.," This corrects for a number of effects such as bias, flat-field and distortions."
 The coronagraphs take sequences of three polarised observations. which are combined to produce total brightness images.," The coronagraphs take sequences of three polarised observations, which are combined to produce total brightness images."
 The COR] and COR2 observations were used to produce standard running difference images., The COR1 and COR2 observations were used to produce standard running difference images.
 The I] observations were background subtracted., The HI observations were background subtracted.
 Modified running dillerence images were then created. which account for the motion ofthe back-ground star field (Maloneyetal. 2009)..," Modified running difference images were then created, which account for the motion ofthe back-ground star field \citep{Maloney:2009p6617}. ."
"A=AL/(1+A2uf/Ac). withAx and Aj evaluated at the outer shock and inner boundary respectively in terms οἱ u,. and u,;. the advection velocities at these locations (see Appendix D for more details).","$\lambda _{\pm}^{\prime}=
\lambda _{\pm}/\left(1+\lambda _{\pm}u/i\omega\right)$, with$\lambda _s$ and $\lambda _i$ evaluated at the outer shock and inner boundary respectively in terms of $u_{rs}$ and $u_{ri}$, the advection velocities at these locations (see Appendix B for more details)."
 Then in equation D6 we approximate and where we have put aad neglect terms of 20345το and higher., Then in equation B6 we approximate and where we have put and neglect terms of $u^2/c_s^2$ and higher.
". These then give. If uv,=0. equation 22 is very similar (o equation 13a of Vishniac&Raw(1989). . ancl is identical if we take r— o. L——L (to agree with their signconvention). /(/+1)/r?—4? "," These then give If $u_r=0$, equation 22 is very similar to equation 13a of \citet{vishniac89}, , and is identical if we take $r\rightarrow\infty$ , $L\rightarrow -L$ (to agree with their signconvention), $l\left(l+1\right)/r^2\rightarrow k^2$ "
The DSSI provides simultaneous observations in two filters by emploving a dichroic beam splitter ancl (wo identical EAICCDs as the imagers.,The DSSI provides simultaneous observations in two filters by employing a dichroic beam splitter and two identical EMCCDs as the imagers.
" We observed Nepler-15 simultaneously in VO and R bandpasses where “V"" has a central wavelength ofAA... and “RO has a central wavelength ofAA... and each filter has aAA.. The details of how we obtain. reduce. and analyze the speckle results and specifics about how they are used to eliminate false positives and aid in (ransit detection are described in llowell et al. ("," We observed Kepler-15 simultaneously in ""V"" and ""R"" bandpasses where ""V"" has a central wavelength of, and ""R"" has a central wavelength of, and each filter has a. The details of how we obtain, reduce, and analyze the speckle results and specifics about how they are used to eliminate false positives and aid in transit detection are described in Howell et al. ("
2011).,2011).
 The speckle observations of the Nepler-15 were obtained on 2010 October 24 (UT) and consisted of [ive sets of 1000. 40 msec individual speckle images.," The speckle observations of the Kepler-15 were obtained on 2010 October 24 (UT) and consisted of five sets of 1000, 40 msec individual speckle images."
 Our R-band reconstructed image is shown in refspeckle with details of the image composition described in Lowell et al. (, Our R-band reconstructed image is shown in \\ref{speckle} with details of the image composition described in Howell et al. (
2011).,2011).
 Along with a nearly identical V-band reconstructed image. (he speckle results reveal no companion star near Ixepler-15 within the annulus from 0.05 to 1.8 arcsec (o a limit of (50) 3.52 magnitudes fainter in R ancl 3.16 magnitudes fainter in. V relative to the Ix=13.76 target star.," Along with a nearly identical V-band reconstructed image, the speckle results reveal no companion star near Kepler-15 within the annulus from 0.05 to 1.8 arcsec to a limit of $\sigma$ ) 3.52 magnitudes fainter in R and 3.16 magnitudes fainter in V relative to the $_{\rm p}=13.76$ target star."
 As a result of the clireet imaging of the field around NKepler-15 we found that two additional stars (besides KIC 11359883) are located wilhin the optimal aperture of IXepler-15., As a result of the direct imaging of the field around Kepler-15 we found that two additional stars (besides KIC 11359883) are located within the optimal aperture of Kepler-15.
" However. both stars are fainter than Ix,=19 and have a negligible effect on the photometry."," However, both stars are fainter than $_{\rm p}=19$ and have a negligible effect on the photometry."
 Onlv KIC 11359883 (IX= 15.99) has a significant effect and we take the diluting elect of its light contribution into account for the light curve modeling., Only KIC 11359883 $_{\rm p}=15.99$ ) has a significant effect and we take the diluting effect of its light contribution into account for the light curve modeling.
 We performed precise RV. follow-up observations of IXepler-15. with the HET (Ramsey el al., We performed precise RV follow-up observations of Kepler-15 with the HET (Ramsey et al.
 1993) and its IRS spectrograph (Tull 1993)., 1998) and its HRS spectrograph (Tull 1998).
 The queue-scheduled observing mode ol the HET usually leads (ο the situation (hat on a given night. data lor many different projects and with different instruments are obtained.," The queue-scheduled observing mode of the HET usually leads to the situation that on a given night, data for many different projects and with different instruments are obtained."
 The observations are ranked according to priorities distributed bv the LET time-allocation-committees as well as additional timing constraints., The observations are ranked according to priorities distributed by the HET time-allocation-committees as well as additional timing constraints.
 We entered. Nepler-15 into the HET queue to be observed in a quasi-rancdom fashion wilh a cadence of a few days to allow proper sampling of the suspected 4.9 d RV orbit., We entered Kepler-15 into the HET queue to be observed in a quasi-random fashion with a cadence of a few days to allow proper sampling of the suspected 4.9 d RV orbit.
 We observed this target from 2010 March 29 until 2010 November 9., We observed this target from 2010 March 29 until 2010 November 9.
 We collected 24 IRS spectra with the I»-cell in the light path for precise RV. measurements., We collected 24 HRS spectra with the $_2$ -cell in the light path for precise RV measurements.
" Furthermore. we obtained one spectrum without the I»-cell to serve as a stellar ""template"" for the RV computation aud to better characterize (he properties of the host star."," Furthermore, we obtained one spectrum without the $_2$ -cell to serve as a stellar “template” for the RV computation and to better characterize the properties of the host star."
 Because of (he faintness of this star. (he ILS setup we emploved for the RV observations," Because of the faintness of this star, the HRS setup we employed for the RV observations"
"ortho-para interconversion of H» takes around 104 yyrs (typical timescale for decrease of o/p-H» ratio by a factor 1/e, see reftime,vol)).","ortho-para interconversion of $_2$ takes around $^4$ yrs (typical timescale for decrease of $_2$ ratio by a factor 1/e, see \\ref{time_evol}) )."
 The free-fall timescale varies as γη and the chemical timescale varies even slower with the density., The free-fall timescale varies as $\sqrt{n}$ and the chemical timescale varies even slower with the density.
" The timescale for o/p-H» conversion is 6x10? yrs at density n 110° cm, and 9x10? and 1.7x10* yrs for densities of 10? and 10* cm, respectively."," The timescale for $_2$ conversion is $\times$ $^3$ yrs at density $n$ $^6$ $^{-3}$, and $\times$ $^3$ and $\times$ $^4$ yrs for densities of $^5$ and $^4$ $^{-3}$, respectively."
" At high densities, the chemical timescales are somewhat shorter than the free-fall timescales (thus roughly one order of magnitude smaller than the empirical core timescales), and the difference even increases at lower densities."," At high densities, the chemical timescales are somewhat shorter than the free-fall timescales (thus roughly one order of magnitude smaller than the empirical core timescales), and the difference even increases at lower densities."
" This comparison between timescales justifies at first order to consider the chemistry independently of the dynamical evolution, as already discussed by?."," This comparison between timescales justifies at first order to consider the chemistry independently of the dynamical evolution, as already discussed by."
". Because the evolutionary stage of the H-MMI core is not constrained, in the following we will focus on the study of the chemical composition when the steady-state is obtained."," Because the evolutionary stage of the H-MM1 core is not constrained, in the following we will focus on the study of the chemical composition when the steady-state is obtained."
 This choice avoids including the age of the core as an additional parameter but will not invalidate the generality of our conclusions., This choice avoids including the age of the core as an additional parameter but will not invalidate the generality of our conclusions.
" As shown in Fig. 7,,"," As shown in Fig. \ref{time_evol},"
 taking the chemical composition at long times overestimates the p-D2H*/o-H2D* ratio in case the core is not yet in chemical steady-state., taking the chemical composition at long times overestimates the $_2$ $^+$ $_2$ $^+$ ratio in case the core is not yet in chemical steady-state.
 This note will be important for the extrapolation of our conclusions in the next sections., This note will be important for the extrapolation of our conclusions in the next sections.
" We could in principle derive the density profile of the core from the continuum emission at 4m. However, this would require the knowledge of the temperature profile, which is not known."," We could in principle derive the density profile of the core from the continuum emission at $\mu$ m. However, this would require the knowledge of the temperature profile, which is not known."
 The Gould Belt Survey made with the Space Observatory should allow to derive accurate temperature profiles., The Gould Belt Survey made with the Space Observatory should allow to derive accurate temperature profiles.
" Until then, as we cannot constrain the density profile without assumptions on the temperature, we restrict ourselves here in running chemical models at different densities, temperatures, and levels of CO depletion."," Until then, as we cannot constrain the density profile without assumptions on the temperature, we restrict ourselves here in running chemical models at different densities, temperatures, and levels of CO depletion."
 Our goal is to investigate under which average conditions the model can reproduce the observations., Our goal is to investigate under which average conditions the model can reproduce the observations.
" A full physical model of the source will be presented in a forthcoming detailed study, after collecting new observational constraints."," A full physical model of the source will be presented in a forthcoming detailed study, after collecting new observational constraints."
 Figure 9 shows the prediction of the model for the D5H*/o-H5D* ratio., Figure \ref{obs_mod} shows the prediction of the model for the $_2$ $^+$ $_2$ $^+$ ratio.
" The thin lines are the model predictions, as a function of temperature, and for different densities and CO depletion factors."," The thin lines are the model predictions, as a function of temperature, and for different densities and CO depletion factors."
" At fixed density, the p-D2H*/o-H2D* ratio is increasing with the CO depletion level, because HD becomes more and more competitive over CO for reaction with"," At fixed density, the $_2$ $^+$ $_2$ $^+$ ratio is increasing with the CO depletion level, because HD becomes more and more competitive over CO for reaction with"
mas (Testi et al. 1998)).,mas (Testi et al. \cite{Tea98}) ).
 We thus decided to perform VLA lich sensitivity radio continuum observations to detect the faint ceutimetric frec-free cussion expected from such an object., We thus decided to perform VLA high sensitivity radio continuum observations to detect the faint centimetric free-free emission expected from such an object.
 The rrevion was observed with the VLA in the period May-June 1998 in the radio coutimuu at 3.6 and 1.2 cin. aud on 26 Jaunary 1999 at 0.7 ane PAG u.," The region was observed with the VLA in the period May-June 1998 in the radio continuum at 3.6 and 1.3 cm, and on 26 January 1999 at 0.7 and 2 cm."
 The observing ouanmeters are suninarized in Table 1.., The observing parameters are summarized in Table \ref{tobs}.
 The 2 cii dataset was obtained as a byproduct of the Lie u experiment: all antennas equipped with Q-baud receivers availabe at the ine of the observaions (12) were used at 0.7 cin while he remaining 15 were eiploved a 2 cu., The 2 cm dataset was obtained as a byproduct of the 0.7 cm experiment: all antennas equipped with Q-band receivers available at the time of the observations (12) were used at 0.7 cm while the remaining 15 were employed at 2 cm.
" At 0. cni we used a fast-switeliiug observing evcle wi havso 1οτος and LO s on-calibraOr (767 away). which resuled in a otal switching cvele of 2160 s aux an cficiency of SHO,"," At 0.7 cm we used a fast-switching observing cycle with 80 s on-source and 40 s on-calibrator $\sim$ $^\circ$ away), which resulted in a total switching cycle of $\sim$ 160 s and an efficiency of $\sim$."
 Uourly pointing sessions on the phase calibrator a 3.6 0 were used to correct for poiutiug ¢rifts at 1.3 aucl 0.7 cia., Hourly pointing sessions on the phase calibrator at 3.6 cm were used to correct for pointing drifts at 1.3 and 0.7 cm.
 3C286 and/or 3C18 were observed to sot the fux scale. which is expected to be accurate withinL5%.," 3C286 and/or 3C48 were observed to set the flux scale, which is expected to be accurate within."
. All data ediins. calibration and inagng Were perforned within the AIPS software package.," All data editing, calibration and imaging were performed within the AIPS software package."
 After standi flux :ud coniplex gain calibration. cach dataset was ποσαανατος using one plase-only aud oue phase and aditud| iteration.," After standard flux and complex gain calibration, each dataset was self-calibrated using one phase-only and one phase and amplitude iteration."
 Consistency along maps at differe freqienceies provided au internal consisteucv check ZI calibration procedures., Consistency among maps at different frequencies provided an internal consistency check of our calibration procedures.
 C'onmparisou of our fluxSs witi previous observations for colpoucut D provided zu adclitiona check., Comparison of our fluxes with previous observations for component D provided an additional check.
 At 0.7 cii heavy data editing Was redqtired «lue o poor atmospheric conditions. ouly Ml of the etire dataset was used to produce the final lips. colTOs]xπας o the last ~30 imiuutes of the run when ΜΕric fluctuatious settled.," At 0.7 cm heavy data editing was required due to poor atmospheric conditions, only $\sim$ of the entire dataset was used to produce the final maps, corresponding to the last $\sim$ 30 minutes of the run when atmospheric fluctuations settled."
 All maps presented here have hee- obtained using the ΑΠ INIACR. task with uniforiiu wegiting of the visibilities and with the ROBUST parameter set to zero., All maps presented here have been obtained using the AIPS IMAGR task with uniform weighting of the visibilities and with the ROBUST parameter set to zero.
 Ta all cases we imaged an area at least equal to the primary beam EWIIM (see Table 13) to se:wel for enissiou., In all cases we imaged an area at least equal to the primary beam FWHM (see Table \ref{tobs}) ) to search for emission.
 No correction for primary beam attenuation has heen applied at any frequency., No correction for primary beam attenuation has been applied at any frequency.
" Tn Figure | we show our radio coutiuuuni images at 3.6. 2.0. 1.3. and 0.7 em of the region coutainius the kuown cuicontiuuuni coniponeuts DE (e.gο, Cosaroui et al. 1991))."," In Figure \ref{fcont} we show our radio continuum images at 3.6, 2.0, 1.3, and 0.7 cm of the region containing the known cm-continuum components D–E (e.g. Cesaroni et al. \cite{Cea94}) )."
 Componcn sA. B. aud € (Caray et al. 1993))," Components A, B, and C (Garay et al. \cite{Gea93}) )"
 are detected in some of OUL aps depending ou scusitivity and (v. 0) coverage. and zwe outside the shown area.," are detected in some of our maps depending on sensitivity and $u,v$ ) coverage, and are outside the shown area."
 Iu addilon to all the previously known ci-coutinuuni conrponeuts. at 3.6 and 1.3 cur we detect four additional SOULCOR. labelled FP to L all above the σ level of ~ Ες at 3.6 cni.," In addition to all the previously known cm-continuum components, at 3.6 and 1.3 cm we detect four additional sources, labelled F to I, all above the $\sigma$ level of $\sim 0.14$ mJy/beam at 3.6 cm."
 At 2.0 cua onv coniponent F is not detected., At 2.0 cm only component F is not detected.
 In Table 2.. 3.6 cin peak positions aud integrated fixes or upper linüts at cach freqirencey for each of the newly detected radio continua couponents are reported: all the newv detected sources are unresolved by our svutliesised beams.," In Table \ref{tspar}, 3.6 cm peak positions and integrated fluxes or upper limits at each frequency for each of the newly detected radio continuum components are reported; all the newly detected sources are unresolved by our synthesised beams."
 A detailed study of all the detected sources eoes bevoud the scope of the preseut etter and will be prescuted in a forheomuneC» paper., A detailed study of all the detected sources goes beyond the scope of the present letter and will be presented in a forthcoming paper.
 In Fieue 2 we show the position of the ΠΟ and ΟΠ masers (Estor Caswell 1989:: Tot:LCL Churelawell 19963). axd of the ΕΕ coniponent F (Hofver et al. 1996))," In Figure \ref{fcontq} we show the position of the $_2$ O and OH masers (Forster Caswell \cite{FC89}; Hofner Churchwell \cite{HC96}) ), and of the mm-continuum component F (Hofner et al. \cite{Hea96}) )"
 aud our 3.6 em (thin cotours) aud he thermal NIT4(5.5) (thick contours. Hofucer et al. ]0013) ," and our 3.6 cm (thin contours) and the thermal $_3$ (5,5) (thick contours, Hofner et al. \cite{Hea94}) )"
Laps cwerlaid on the 2.2 san near infrared iuage from Testi ο al. (1998))., maps overlaid on the 2.2 $\mu$ m near infrared image from Testi et al. \cite{Tea98}) ).
 Within the astrometric uicertainties (κ ‘for t16 NIB. data. <0.2” Tor all the other data). the newly discovered. curconti Inτος called Ε in Table 2 is coincident with 10 3(5.5rFr) IIC. the niuui-coutiunmua ar the NIR source.," Within the astrometric uncertainties $\le$ for the NIR data, $\le$ for all the other data), the newly discovered cm-continuum source called F in Table \ref{tspar} is coincident with the $_3$ (5,5) HC, the mm-continuum and the NIR source."
 T IC aid nmumuconponeut F locaed between the Clu-CCinuuni sources D and E was the primary target of our observatiois., The HC and mm-component F located between the cm-continuum sources D and E was the primary target of our observations.
 We detected source F a 3.6 and 1.3 ci axd set upper Iunits at 2.0 and 0.7 ci., We detected source F at 3.6 and 1.3 cm and set upper limits at 2.0 and 0.7 cm.
 In Figure 3) we show tie radio contimuni spectruni of source F. Data from he preseit work are preseuted as filed circles. while the on circle is from C'esaroui et al. (199 1).," In Figure \ref{fspec} we show the radio continuum spectrum of source F. Data from the present work are presented as filled circles, while the open circle is from Cesaroni et al. \cite{Cea94}) ),"
 aud the open scuare from Iofuer et al. (1996)., and the open square from Hofner et al. \cite{Hea96}) ).
"Throughout this paper. we assume the following cosmologv li,=71kmsAlpe Ox,=027 and O4—70.73. ina flat Universe.","Throughout this paper, we assume the following cosmology: $H_{0} =
71\; {\rm km \; s^{-1} \; Mpc^{-1}}$ , $\Omega_{\rm M} = 0.27$ and $\Omega_{\rm \Lambda} = 0.73$, in a flat Universe."
 At the redshift of the target 1 aresce = 5.007 kpc., At the redshift of the target 1 arcsec = 5.007 kpc.
 The spectral index à is defined as ο) xp 10510-080. was target. of Space-VLDBI observations., The spectral index $\alpha$ is defined as $S$ $\nu$ ) $\propto \nu^{- \alpha }$ 1510-089 was target of Space-VLBI observations.
 I was observed at 4.8 Gllz (€ band) by VLBA|ILALC'A on 1999 August 11 and on 2000 Alay 13 in single polarization mode., It was observed at 4.8 GHz (C band) by VLBA+HALCA on 1999 August 11 and on 2000 May 13 in single polarization mode.
 In each observing run the target was observed. for about 6 hours., In each observing run the target was observed for about 6 hours.
 Examples of thewe coverage and amplitude vs baseline length are presented in Fig. 1.., Examples of the coverage and amplitude vs baseline length are presented in Fig. \ref{uv-vsop}.
 The target source. was observed with the VLBA with the aclelition of the Efelsberg telescope at sA Cillz (X band) ancl 22.2 11 Osx band). in full polarization mode with a bandwidth of 32 Mllz at 128. Mbps.," The target source was observed with the VLBA with the addition of the Effelsberg telescope at 8.4 GHz (X band) and 22.2 GHz (K band), in full polarization mode with a bandwidth of 32 MHz at 128 Mbps."
 To study possible changes in the source structure. subsequent VLBA observations at SA Cllz were carried. out in full polarization with a recording bandwidth of 32 Mllz at 128 Mbps.," To study possible changes in the source structure, subsequent VLBA observations at 8.4 GHz were carried out in full polarization with a recording bandwidth of 32 MHz at 128 Mbps."
 The correlation was performed. at the. VLBA correlator in Socorro and the data reduction was carried out with the NRAO ALPS package., The correlation was performed at the VLBA correlator in Socorro and the data reduction was carried out with the NRAO AIPS package.
 In the XN. band the instrumental polarization was removed by using the ALPS task PCAL., In the X band the instrumental polarization was removed by using the AIPS task PCAL.
 The absolute orientation of the electric vector of the calibrator B1I749|096 was then compared with the VLA/VLDA polarization calibration database to derive the corrections., The absolute orientation of the electric vector of the calibrator B1749+096 was then compared with the VLA/VLBA polarization calibration database to derive the corrections.
 Phe values derived are in good agreement. {5 57)., The values derived are in good agreement $\leq$ $^{\circ}$ ).
 Final images for each epoch and at each [requeney were produced after a number of phase. self-calibration iterations., Final images for each epoch and at each frequency were produced after a number of phase self-calibration iterations.
 Aniplituce self-calibration was applied. at the enc of the process using a solution interval longer than the scan length. to remove residual svstematic errors and to fine tune the [lux density scale. but not to force the individual data points to follow the mocel.," Amplitude self-calibration was applied at the end of the process using a solution interval longer than the scan length, to remove residual systematic errors and to fine tune the flux density scale, but not to force the individual data points to follow the model."
 The resolution ab the various frequencies are almost. comparable CLable 1) due to we similar wreecoverage reached. with the different interferometers used., The resolution at the various frequencies are almost comparable (Table \ref{vlba_obs}) ) due to the similar -coverage reached with the different interferometers used.
 At S4 CGllIz. besides the total intensity (D). images in the Stokes’ U and. (Q parameters were produced with the final fully calibrated datasets.," At 8.4 GHz, besides the total intensity (I), images in the Stokes' U and Q parameters were produced with the final fully calibrated datasets."
 Final VLBI images at 4.8. 8.4 and 22.2 Gllz are shown in Fig. 2..," Final VLBI images at 4.8, 8.4 and 22.2 GHz are shown in Fig. \ref{vlba_images}."
 For the observations carried out on 1999 January 11. we produced also à low-resolution image at 22.2 (111 using the same range. restoring beam and image sampling of the SA Cz data in order to produce the spectral index image. which is presented in Fig.," For the observations carried out on 1999 January 11, we produced also a low-resolution image at 22.2 GHz using the same range, restoring beam and image sampling of the 8.4 GHz data in order to produce the spectral index image, which is presented in Fig."
 3. superimposed on the 22.2 Cillz contours obtained from the low-resolution image., \ref{spix} superimposed on the 22.2 GHz contours obtained from the low-resolution image.
 Information on the VLBI observations is reported in Table l.., Information on the VLBI observations is reported in Table \ref{vlba_obs}.
 To study the radio variability. of LI510-089. we compared. our observations with multi-epoch 15-Cillz (U band) VLBA cata from the MO.LAVIS programme spanning a time interval from 1995. July 28 to 2010. December 23.," To study the radio variability of 1510-089, we compared our observations with multi-epoch 15-GHz (U band) VLBA data from the MOJAVE programme spanning a time interval from 1995 July 28 to 2010 December 23."
 The typical resolution is about 0.5 mas., The typical resolution is about $\times$ 0.5 mas.
 For each of the 5] epochs analysed. we imported the calibrated uie-datasets (Listeretal.2009¢) into the NRAO ALPS package and performed a [ον phase-only selt-calibration iterations before producing the final total intensity images which resulted to be fully consistent with those reported. by Listeretal. (2009c).," For each of the 51 epochs analysed, we imported the calibrated -datasets \citep{lister09a} into the NRAO AIPS package and performed a few phase-only self-calibration iterations before producing the final total intensity images which resulted to be fully consistent with those reported by \citet{lister09a}."
. For those datasets in full polarization mode. we produced also Stokes’ U and Q images.," For those datasets in full polarization mode, we produced also Stokes' U and Q images."
 Ehe rms noise level measurecl on the image plane is in the range of 0.15 and 0.3 nily/beam., The rms noise level measured on the image plane is in the range of 0.15 and 0.3 mJy/beam.
 Total intensity images concerning the observing epochs between July 1995 and July 2008 are published. in Listeretal.(2009c)., Total intensity images concerning the observing epochs between July 1995 and July 2008 are published in \citet{lister09a}.
. An image of the source structure in September 2010 is presented in Fig., An image of the source structure in September 2010 is presented in Fig.
 4. as an The lux density and deconvolved. angular size of cach sourcecomponent were measured by means of the ALPS task JMEUL. which performs a Gaussian [it to the source," \ref{mojave_sep10} as an The flux density and deconvolved angular size of each sourcecomponent were measured by means of the AIPS task JMFIT, which performs a Gaussian fit to the source"
defined to increase counter-clockwise) are significantly higher than along the galactic minor axis at the same radius.,defined to increase counter-clockwise) are significantly higher than along the galactic minor axis at the same radius.
" This is strong evidence that the halo profile is not spherically-symmetric, but is rather flattened in the direction of the disk."," This is strong evidence that the halo profile is not spherically-symmetric, but is rather flattened in the direction of the disk."
" Another possibility is that we are not seeing a true stellar halo, but rather a very extended thick disk, which would necessarily have a very flattened distribution (see discussion in Section ?? below)."," Another possibility is that we are not seeing a true stellar halo, but rather a very extended thick disk, which would necessarily have a very flattened distribution (see discussion in Section \ref{sec:vertprofile} below)."
 The azimuthal trend of the halo density in Section 4.1.0 strongly indicates that the stellar halo of NGC 253 is flattened., The azimuthal trend of the halo density in Section \ref{sec:radprofile} strongly indicates that the stellar halo of NGC 253 is flattened.
" Therefore, in order to determine an accurate radial profile, we must adopt “annuli” that mimic the intrinsic distribution of the stars."," Therefore, in order to determine an accurate radial profile, we must adopt “annuli” that mimic the intrinsic distribution of the stars."
" In this section, we adopt elliptical annuli aligned with the stellar disk."," In this section, we adopt elliptical annuli aligned with the stellar disk."
" For the elliptical axis ratio, we have tested values ranging from the observed disk axis ratio of b/a=0.25 (??),, up to b/a=0.5."," For the elliptical axis ratio, we have tested values ranging from the observed disk axis ratio of $b/a=0.25$ \citep{pence80,koribalski-etal04}, up to $b/a=0.5$."
" In the azimuthal direction, we adopt an angular coordinate @ defined by where Tmaj and rjj, are the projections of the radial vector along the major and minor axes respectively."," In the azimuthal direction, we adopt an angular coordinate $\phi$ defined by where $r_{\mathrm{maj}}$ and $r_{\mathrm{min}}$ are the projections of the radial vector along the major and minor axes respectively."
" If the halo distribution were perfectly disk-like, angular sectors spaced uniformly in ϕ would correspond to equal volumes."," If the halo distribution were perfectly disk-like, angular sectors spaced uniformly in $\phi$ would correspond to equal volumes."
 Our results are qualitatively unchanged if we instead adopt sectors that are spaced uniformly in the true projected azimuth 0., Our results are qualitatively unchanged if we instead adopt sectors that are spaced uniformly in the true projected azimuth $\theta$.
" The elliptical annuli had minor axes spaced by 4 kpc, which were further separated into sectors with azimuthal extent Ad=8°, as shown in Figure 7 (Aó values from 4? to 10° showed similar behaviour)."," The elliptical annuli had minor axes spaced by $4$ kpc, which were further separated into sectors with azimuthal extent $\Delta\phi = 8\degr$, as shown in Figure \ref{fig:ellipsectors}
 $\Delta\phi$ values from $4\degr$ to $10\degr$ showed similar behaviour)."
 The radial halo profile in elliptical annuli with axis ratios b/a—0.35 is shown in the main panel of Figure 7.., The radial halo profile in elliptical annuli with axis ratios $b/a=0.35$ is shown in the main panel of Figure \ref{fig:ellipsectors}.
 A single power-law with a slope of —2.8+0.3 provides an acceptable fit to the data., A single power-law with a slope of $-2.8 \pm 0.3$ provides an acceptable fit to the data.
 This further demonstrates that the break in the power law seen in Figure 6 is entirely an artifact of the assumed geometry: the circular annuli coupled with the 5 kpc inner cutoff probe different ranges of elliptical annuli as a function of radius., This further demonstrates that the break in the power law seen in Figure \ref{fig:radsectors} is entirely an artifact of the assumed geometry: the circular annuli coupled with the $5$ kpc inner cutoff probe different ranges of elliptical annuli as a function of radius.
" Shifting the inner points up by 0.2 dex, as in Section 4.1.0,, steepens the slope to n——3.4."," Shifting the inner points up by $0.2$ dex, as in Section \ref{sec:radprofile}, steepens the slope to $n=-3.4$."
 The azimuthal profiles are shown in the bottom-left inset of Figure 7.., The azimuthal profiles are shown in the bottom-left inset of Figure \ref{fig:ellipsectors}.
" Unlike with the circular annuli, the densities do not show a systematic gradient, again demonstrating the flattening of the halo."," Unlike with the circular annuli, the densities do not show a systematic gradient, again demonstrating the flattening of the halo."
" However, the azimuthal profiles are far from flat, and exhibit a strong"," However, the azimuthal profiles are far from flat, and exhibit a strong"
right.,.
".. in the last cooling case 5i;<rjy. Where is (he observed peak fIux at the Iuminosity distance D=(12:)IM""cz)dl/dz|dz with dl/dz!—(1+-Mu[OuDy!Qj2 aad = [B2amype, he. is the comoving magnetic field."," in the fast cooling case $\nu_{c,f}<\nu_{m,f}$, where is the observed peak flux at the luminosity distance $D=(1+z) \int_{0}^{z} (1+z) |dt/dz| dz$ with $|dt/dz|^{-1}=(1+z) H_{0} [\Omega_{m}(1+z)^{3}+\Omega_{\Lambda}]^{1/2}$, and =[32 n c, is the comoving magnetic field."
" The characteristic svuchrotron requency 5,5; and the cooling requency ÉL.y are given by 7(5,,5) and r(5,y). respectively. where is the frequeney al which electrons with the Lorentz factor +, radiate. ancl 0001)Note that we may use By and 5,,,; in equations (À21)) and (A23)) even in the phase."," The characteristic synchrotron frequency $\nu_{m,f}$ and the cooling frequency $\nu_{c,f}$ are given by $\nu(\gamma_{m,f})$ and $\nu(\gamma_{c,f})$, respectively, where is the frequency at which electrons with the Lorentz factor $\gamma_{e}$ radiate, and -1),Note that we may use $B_{f}$ and $\gamma_{m,f}$ in equations \ref{eq:Bf}) ) and \ref{eq:gmf}) ) even in the non-relativistic phase."
 The svnchrotron self-absorption limits the flix below the blackbody emission with the electron temperature (e.g..Sari&Piran 1999)..," The synchrotron self-absorption limits the flux below the blackbody emission with the electron temperature \citep[e.g.,][]{sari99a}. ."
 This is given bv, This is given by
auc dissipation processes in tle pulsar magnetosphere. since the maximum enerey 1s determined by he acceleratiou-raciation-reaction-limit for typical energetic pulsars.,"and dissipation processes in the pulsar magnetosphere, since the maximum energy is determined by the acceleration-radiation-reaction-limit for typical energetic pulsars."
 The 5-ray emission region has herefore been explored by comparing theoretical models with the observed light curve(e.g.. WattersVenter.Harding&Cuillemot(2009):RomaniWatters (2010))).," The $\gamma$ -ray emission region has therefore been explored by comparing theoretical models with the observed light curve(e.g., \citet{Wa09, VHG09, RW10}) )."
 The pulsed emissiou is also detected in other energy bauds (X-ray. ultraviolet. optical aud radio) for some sources (e.g.. Thompson (2001))).," The pulsed emission is also detected in other energy bands (X-ray, ultraviolet, optical and radio) for some sources (e.g., \citet{T04}) )."
 The spectral features are non-thermal except for the soft. N-ray rauge. aud he Leht curves [rom a siugle object are. in general. dillerent [rom oue energy. band to another.," The spectral features are non-thermal except for the soft X-ray range, and the light curves from a single object are, in general, different from one energy band to another."
 For example. profiles of the light curve in one spin period are clilferent iu the ~- aud X-ray ranges in the Vela pulsar (Abdoetal.2009a)..," For example, profiles of the light curve in one spin period are different in the $\gamma$ - and X-ray ranges in the Vela pulsar \citep{Abvela2}."
 The peak phase of dillerent οποίον range is expected to coincide. siuce the emittingi particles are related to a pair cascade process.," The peak phase of different energy range is expected to coincide, since the emitting particles are related to a pair cascade process."
 However. tlie observation shows that the pliase clepeuds ou the enerey bauds.," However, the observation shows that the phase depends on the energy bands."
 This meaus that their emission regions are wot the same region., This means that their emission regions are not the same region.
 A complete understanding of light curve behavior in multi-waveleneth bauds can provide valuable iuformatiou about the particle acceleration region., A complete understanding of light curve behavior in multi-wavelength bands can provide valuable information about the particle acceleration region.
 Note that we do uot discuss soft X-rays. which are believed to be thermal radiatiou from the —neutron star surface (e.g.. (2002))).," Note that we do not discuss soft X-rays, which are believed to be thermal radiation from the neutron star surface (e.g., \citet{Ja02}) )."
 Possible origius of nou-tlieruial pulsed emissions-- have been considered in the polar cap 1996).. slot. eap (Muslimov&Hardinge22001).. and outer gap (Cheug. —odels.," Possible origins of non-thermal pulsed emissions have been considered in the polar cap \citep{DH96}, slot gap \citep{MH04}, and outer gap \citep{CHR86} models."
 RecentFerinz observatious with high 5-ray. photon nunber statistics have showed that the phase-averagecd spectrum above 200 MeV is well fitted by a power law plus exponential cutoll. aud that a cutolf shape sharper than a simple exponential is rejected. with high signilicance (e.e.. Abdoetal. (2009a))).," Recent observations with high $\gamma$ -ray photon number statistics have showed that the phase-averaged spectrum above 200 MeV is well fitted by a power law plus exponential cutoff, and that a cutoff shape sharper than a simple exponential is rejected with high significance (e.g., \citet{Abvela2}) )."
 This rules out the near-surface emission. proposed iu polar cap cascade models (Daugherty&Harding 1996).. which would exhibit a much sharp spectral cutolldue to iaguetie pair-productiou attenuation.," This rules out the near-surface emission proposed in polar cap cascade models \citep{DH96}, , which would exhibit a much sharp spectral cutoff due to magnetic pair-production attenuation."
 Thus. pulsed οταν emission originates in the outer magnetospliere. as considered in the outer gap moclel.," Thus, pulsed $\gamma$ -ray emission originates in the outer magnetosphere, as considered in the outer gap model."
 Takata.Chang&Shibata(2008) (hereafter TCS05) cousidered a tlree-dimensioual geometrical euission model to fit the observed light curves at differeut euerey bauds., \citet{TCS08} (hereafter TCS08) considered a three-dimensional geometrical emission model to fit the observed light curves at different energy bands.
 The model is exteudec with some model parameters [rom the calculated gap structure in a two-dimensioual merician plane., The model is extended with some model parameters from the calculated gap structure in a two-dimensional meridian plane.
 By comparing the light curves of tlie Vela pulsar. they fouud that the X-ray. emission is produced by both imward and outward emission from the gap region. and that UV/optical emission originates [rom secondary pairs at a higlier altitude.," By comparing the light curves of the Vela pulsar, they found that the X-ray emission is produced by both inward and outward emission from the gap region, and that UV/optical emission originates from secondary pairs at a higher altitude."
 The nuuber of light curves of pulsars observed at >-ray aud other energy bands is increasing thanks toFern so that it is worthwhile to investigate whether outer gap model is applicable to other sources.," The number of light curves of pulsars observed at $\gamma$ -ray and other energy bands is increasing thanks to, so that it is worthwhile to investigate whether outer gap model is applicable to other sources."
 In this paper. we investigate the emission regious of several pulsars by fitting the simplified moclel of TCS08 to the observed multi-wavelength liglit curves.," In this paper, we investigate the emission regions of several pulsars by fitting the simplified model of TCS08 to the observed multi-wavelength light curves."
 In this model. we have to specify the locatious of the upper aud lower boundaries of the gap region where the nou-corotation potential is zero.," In this model, we have to specify the locations of the upper and lower boundaries of the gap region where the non-corotation potential is zero."
 Therefore. we explicitly iutroduce the altitude of the gap region as a parameter. in order to fit the observational data easily.," Therefore, we explicitly introduce the altitude of the gap region as a parameter, in order to fit the observational data easily."
 The light curves also depend ou the dipole incliuation and viewing augles., The light curves also depend on the dipole inclination and viewing angles.
 Iu our metliod. such parameters are elimiuated by other observational data. aid ouly the altitude ischauged for the fitting.," In our method, such parameters are eliminated by other observational data, and only the altitude ischanged for the fitting."
 In the most studies. the lower boundary of the emission regiou," In the most studies, the lower boundary of the emission region"
"for their presence (e.g.,thecompaniontoyoungpulsarPSRJ19064-0746; B006).","for their presence \citep[e.g., the companion to the young pulsar PSR J1906$+$0746;][]{lorimer06}."
". Given that the NS blackbody emission is gravitationally bent, allowing us to view >75% of the NS surfaces in X rays (Beloborodov]2002)., sufficiently deep X-ray observations are virtually guaranteed to detect these MSPs."," Given that the NS blackbody emission is gravitationally bent, allowing us to view $>75\%$ of the NS surfaces in X rays \citep{belo02}, sufficiently deep X-ray observations are virtually guaranteed to detect these MSPs."
" found no correlation between the X-ray and radio luminosities of MSPs in 47 Tuc, as expected due to the differing nature and spatial location of the X-ray and radio emission, and found that all MSPs in 47 have X-ray luminosities ranging betweenπο Lx(0.5—6 keV)=2x10?? and 2x10?! erg s"," \citet{Heinke05a} found no correlation between the X-ray and radio luminosities of MSPs in 47 Tuc, as expected due to the differing nature and spatial location of the X-ray and radio emission, and found that all MSPs in 47 have X-ray luminosities ranging between $L_X (0.5-6$ $) =2\times10^{30}$ and $2\times10^{31}$ erg $^{-1}$."
" showed that the X-ray spectra of the MSPs in 47 Tuc are typically dominated by thermal blackbody-like emission from the NS surface around the polar caps, with temperature 1—3x106 K. This X-ray emission is sometimes overwhelmed by additional non-thermal X rays that are either magnetospheric or due to an intra-binary shock."," \citet{Bogdanov06} showed that the X-ray spectra of the MSPs in 47 Tuc are typically dominated by thermal blackbody-like emission from the NS surface around the polar caps, with temperature $1-3\times10^6$ K. This X-ray emission is sometimes overwhelmed by additional non-thermal X rays that are either magnetospheric or due to an intra-binary shock."
 also showed that there are no clear systematic differences between the X-ray properties of MSPs in 47 Tuc and in the Galactic field., \citet{Bogdanov06} also showed that there are no clear systematic differences between the X-ray properties of MSPs in 47 Tuc and in the Galactic field.
" Thus, we requested anXMM observation capable of detecting any of the known MSPs in 47 Tuc, were they located at the distance of J0917--4638."," Thus, we requested an observation capable of detecting any of the known MSPs in 47 Tuc, were they located at the distance of $+$ 4638."
" J0917+4638 had not previously been observed inthe X-ray since theROSAT All-Sky Survey where it was not detected (unsurprisingly, considering 2000),the short exposure time)."," $+$ 4638 had not previously been observed inthe X-ray since the All-Sky Survey \citep{voges99,fsc}, where it was not detected (unsurprisingly, considering the short exposure time)."
" We observed J09174-4638 on 2008 May 7 for 17 ks (ObsID 0553440101) with XMM's EPIC camera, consisting of two MOS CCD detectors (Turneretαἱ. and a pn CCD detector (Strüderet "," We observed $+$ 4638 on 2008 May 7 for 17 ks (ObsID 0553440101) with 's EPIC camera, consisting of two MOS CCD detectors \citep{turner01} and a pn CCD detector \citep{struder01}. ."
"All data were reduced using and SAS al.[2001)..version We excluded times of soft proton background flaring, when the pn camera's count rate exceeded 25 (0.2—10 keV) counts s!, or when the MOS cameras exceeded 7 or 8 —10 keV) counts s! for the MOS1 or MOS2 cameras (0.2respectively."," All data were reduced using and SAS version We excluded times of soft proton background flaring, when the pn camera's count rate exceeded 25 $0.2-10$ keV) counts $^{-1}$, or when the MOS cameras exceeded 7 or 8 $0.2-10$ keV) counts $^{-1}$ for the MOS1 or MOS2 cameras respectively."
" This left 8.9 ks of good data from the pn detectors, and 11.3 ks from the MOS detectors."," This left 8.9 ks of good data from the pn detectors, and 11.3 ks from the MOS detectors."
" We filtered the events on pixel patterns (trying PATTERN<= or <=4 for pn and for MOS data), and1 for FLAG==0."," We filtered the events on pixel patterns (trying $<=$ 1 or $<=$ 4 for pn and $<=$ 12 for MOS data), and for FLAG==0."
 We PATTERNchoose an «—12energy range of 0.2—1.5 keV to obtain optimal sensitivity to the soft blackbody emission expected from MSPs., We choose an energy range of $0.2-1.5$ keV to obtain optimal sensitivity to the soft blackbody emission expected from MSPs.
 No source is detected at or within 1' of the location of J0917--4638 eitherwith detection algorithms or by eye., No source is detected at or within $1\amin$ of the location of $+$ 4638 eitherwith detection algorithms or by eye.
" We utilize our knowledge of the XMM point spread and absolute pointing accuracy («1""; to determine an upper limit.", We utilize our knowledge of the point spread and absolute pointing accuracy \citep[$<1\asec$;][]{Kirsch04} to determine an upper limit.
" For the pn, 50% of 1.5 keV photons are found within 8"", and 8096 within 20""."," For the pn, $50\%$ of 1.5 keV photons are found within $8\asec$, and $80\%$ within $20\asec$."
" For the MOS cameras, 5096 of 1.5 keV photons are recorded within 8"", and 7596 within 20""."," For the MOS cameras, $50\%$ of 1.5 keV photons are recorded within $8\asec$, and $75\%$ within $20\asec$."
" We find 3 counts within an 8"" circle, or 20 counts within a 20"" circle, in the combined image."," We find 3 counts within an $8\asec$ circle, or 20 counts within a $20\asec$ circle, in the combined image."
" This is consistent with a nondetection, as the expected background counts in these circles are 3.2+0.2 and 19.8+1.3 counts, respectively, as derived from nearby background regions."," This is consistent with a nondetection, as the expected background counts in these circles are $3.2\pm0.2$ and $19.8\pm1.3$ counts, respectively, as derived from nearby background regions."
" We use the X-ray faintest MSP in 47 Tuc, 47 Tuc T, to calibrate our expectations forthe detection of an MSP in J0917--4638."," We use the X-ray faintest MSP in 47 Tuc, 47 Tuc T, to calibrate our expectations forthe detection of an MSP in $+$ 4638."
 47 Tuc T has Lx(0.2—1.5 keV)=1.5x10°° erg s! and a 134 eV blackbody spectrum etal.|2006)., 47 Tuc T has $L_X (0.2-1.5$ $)=1.5\times10^{30}$ erg $^{-1}$ and a 134 eV blackbody spectrum \citep{Bogdanov06}.
". We use to determine the expected EPIC count rates from 47 Tuc T were it located at 2.3 kpc (the distance to J0917--4638)(Dickey behind&Lodanan| an estimated[950), NyC15x10! onc", We use to determine the expected EPIC count rates from 47 Tuc T were it located at 2.3 kpc (the distance to $+$ 4638) behind an estimated $N_H=1.5\times10^{20}$ $^{-2}$ \citep{dickey90}.
" We expect 10.9 counts within 8”, or 17.0 within 20”, accounting for the relevant encircled energy fractions, from such an MSP."," We expect 10.9 counts within $8\asec$, or 17.0 within $20\asec$, accounting for the relevant encircled energy fractions, from such an MSP."
" Comparing the predicted counts with the Poisson errors on the detected counts equation7),, we find that the number of counts within 20” is 3.00 below expectations for the faintest known MSP in 47 Tuc, while the counts within 8” are 3.50 below those expectations."," Comparing the predicted counts with the Poisson errors on the detected counts \citep[][equation 7]{Gehrels86}, we find that the number of counts within $20\asec$ is $3.0\sigma$ below expectations for the faintest known MSP in 47 Tuc, while the counts within $8\asec$ are $3.5\sigma$ below those expectations."
" 47 Tuc T is the X-ray faintest of the 15 independently measured MSPs in 47 Tuc; the median X-ray luminosity is 2.1x greater etαἱ/2006)., which is ruled out at 6.30 confidence."," 47 Tuc T is the X-ray faintest of the 15 independently measured MSPs in 47 Tuc; the median X-ray luminosity is $2.1\times$ greater \citep{Bogdanov06}, which is ruled out at $6.3\sigma$ confidence."
(Bogdanov Our nondetection is therefore strong evidence against the existence of an MSP in J0917+4638., Our nondetection is therefore strong evidence against the existence of an MSP in $+$ 4638.
 We have searched for evidence of an MSP companion to the LMWD J0917+4638 through radio and X-ray observations., We have searched for evidence of an MSP companion to the LMWD $+$ 4638 through radio and X-ray observations.
" Our radio search reaches a sensitivity sufficient to detect roughly two-thirds of the known MSPs, while our X-ray search is sensitive enough to detect any of the 15 independently identified MSPs in 47 Tuc."," Our radio search reaches a sensitivity sufficient to detect roughly two-thirds of the known MSPs, while our X-ray search is sensitive enough to detect any of the 15 independently identified MSPs in 47 Tuc."
" Together, our nondetections provide strong evidence against the presence of an MSP in this system."," Together, our nondetections provide strong evidence against the presence of an MSP in this system."
" Furthermore, since any NS companion to J0917--4638 would presumably have been recycled through accretion from the LMWD, we rule out the presence of a NS in this system."," Furthermore, since any NS companion to $+$ 4638 would presumably have been recycled through accretion from the LMWD, we rule out the presence of a NS in this system."
" Although a black hole companion is still conceivable (as the 7.6 hr orbital period would not induce current accretion and X-ray activity), such a companion is far less probable than a WD companion given both the system's mass function and the stellar initial mass function for M>1Mo (e.g.,"," Although a black hole companion is still conceivable (as the 7.6 hr orbital period would not induce current accretion and X-ray activity), such a companion is far less probable than a WD companion given both the system's mass function \citep{kilic07b} and the stellar initial mass function for $M \geq 1 \ M_\odot$ \citep[e.g.,][]{scalo98}."
 We conclude that J0917--4638's more massive [1998)..companion (M>0.28 Mo) is almost certainly another WD., We conclude that $+$ 4638's more massive companion $M \geq 0.28\ M_\odot$ ) is almost certainly another WD.
 Roughly two dozen binaries are known and in ten such systems both WD WD/WDmasses have been measured (seeandreferences therein).., Roughly two dozen WD/WD binaries are known and in ten such systems both WD masses have been measured \citep[see][and references therein]{nelemans05}. .
 The individualmasses of WDs in these systems range between 0.29 and 0.71 Mo; the medianmass for those with measured masses (and not just lower limits) is 0.43 Mo., The individualmasses of WDs in these systems range between $0.29$ and $0.71\ M_\odot$ ; the medianmass for those with measured masses (and not just lower limits) is $0.43\ M_\odot$ .
" The majority of these double WD systems have mass ratios near unity, which is contrary to what is expected from standard population synthesis models 2005).."," The majority of these double WD systems have mass ratios near unity, which is contrary to what is expected from standard population synthesis models \citep{nelemans05b}. ."
" This has been used to argue that the standard prescription for energycommon balanceenvelope (a-formalism),evolution, should be replaced by"," This has been used to argue that energy balance $\alpha$ -formalism), the standard prescription for common envelope evolution, should be replaced by"
"where for a circular survey with radius ©. Xp m) assundue that -ns(uthe true potential ""ραaverages to zero.","where for a circular survey with radius $\Theta$, = 4 ) - ( 1 - ) + _x + ), assuming that the true potential averages to zero."
 Seence(tera) js an estimate of the deusity field.," Hence ) = _r r^2 _r ) - ) is an estimate of the Newtonian gravitational potential, while ) = ^2 _r r^2 _r )) - ) is an estimate of the density field."
 The derivation of equations (3.1)) and (3.1)) are the main results of this paper. aud demoustrate that the full. 3-D Newtomian potential aud density fields can be reconstructed frou weak leusiug observations.," The derivation of equations \ref{unbiaspot}) ) and \ref{unbiasdel}) ) are the main results of this paper, and demonstrate that the full, 3-D Newtonian potential and density fields can be reconstructed from weak lensing observations."
 In practice \ will uot be a major problem for large surveys where the moments of the true lensing potential will average to zero., In practice $\chi$ will not be a major problem for large surveys where the moments of the true lensing potential will average to zero.
 However for snall fields this way not be so true. aud one may instead wish to set the edge of the field to o=0.," However for small fields this may not be so true, and one may instead wish to set the edge of the field to $\phi=0$."
 Having set out the formal method for a 3-D recoustruction of the dark matter distribution. we now test it on simulations.," Having set out the formal method for a 3-D reconstruction of the dark matter distribution, we now test it on simulations."
 Our simulation consists of a double cluster with a Navarro-Frenk-White profile (Figure 3))., Our simulation consists of a double cluster with a Navarro-Frenk-White profile (Figure \ref{fig:grav1}) ).
 Further details can be found in Bacon aud Tavlor (2003)., Further details can be found in Bacon and Taylor (2003).
 First we examine the accuracy of our mcthod when the shear field is perfectly known evervwhere., First we examine the accuracy of our method when the shear field is perfectly known everywhere.
 Figure l displavs the recoustructed ® aud the difference between iuput aud reconstructed fields., Figure \ref{fig:grav2} displays the reconstructed $\Phi$ and the difference between input and reconstructed fields.
 It cau be secu, It can be seen
"With the attribution of the 11.0 wm feature to PAH* and the 11.2 um to PAH?, we can investigate the possibility of using this ratio to probe the ionization fraction of PAHs in the PDR.","With the attribution of the 11.0 $\mu$ m feature to $^+$ and the 11.2 $\mu$ m to $^0$, we can investigate the possibility of using this ratio to probe the ionization fraction of PAHs in the PDR."
 One of the classic methods to trace the PAH ionization fraction is the [6.2]/[11.3] um integrated intensity ratio (e.g. ?))., One of the classic methods to trace the PAH ionization fraction is the [6.2]/[11.3] $\mu$ m integrated intensity ratio (e.g. \cite{galliano}) ).
" There are other tracers of ionization such as the [7.7]/[11.3] wm and [8.6]/[11.3] uum ratios, but the 6.2, 7.7, and 8.6 um features include blended PAH* and PAH? bands."," There are other tracers of ionization such as the [7.7]/[11.3] $\mu$ m and [8.6]/[11.3] $\mu$ m ratios, but the 6.2, 7.7, and 8.6 $\mu$ m features include blended $^+$ and $^0$ bands."
" As we show here, the 11.0 um band is a purely cationic band and the 11.2 um band is purely neutral, increasing the accuracy of ionization fraction measurements."," As we show here, the 11.0 $\mu$ m band is a purely cationic band and the 11.2 $\mu$ m band is purely neutral, increasing the accuracy of ionization fraction measurements."
" To demonstrate the reliability of the [11.0]/[11.2] um ratio as an ionization indicator, we compare the [6.2]/[11.2] um ratio to the [11.0]/[11.2] jum ratio using the IRS-SL and IRS-SH observations of the NGC 7023-NW (Figure 10))."," To demonstrate the reliability of the [11.0]/[11.2] $\mu$ m ratio as an ionization indicator, we compare the [6.2]/[11.2] $\mu$ m ratio to the [11.0]/[11.2] $\mu$ m ratio using the IRS-SL and IRS-SH observations of the NGC 7023-NW (Figure \ref{fig:linear}) )."
" In order to have an accurate measurement of the 11.2 um feature, without contamination from the 11.0 um satellite feature, we compare the integrated intensity of the 6.2 um feature from IRS-SL observations to the intensity of the 11.2 um feature using the high-resolution observations, since the 11.0 was not resolved and separated in the IRS-SL observations."," In order to have an accurate measurement of the 11.2 $\mu$ m feature, without contamination from the 11.0 $\mu$ m satellite feature, we compare the integrated intensity of the 6.2 $\mu$ m feature from IRS-SL observations to the intensity of the 11.2 $\mu$ m feature using the high-resolution observations, since the 11.0 was not resolved and separated in the IRS-SL observations."
 The maps were re-gridded using so that each point of the SH map corresponds to the same spatial position on the SL map., The maps were re-gridded using so that each point of the SH map corresponds to the same spatial position on the SL map.
 Only the highest signal to noise data were used in this plot., Only the highest signal to noise data were used in this plot.
" For the 6.2 um low-resolution map, we set a band integrated intensity threshold of 10 Wm~*sr7!."," For the 6.2 $\mu$ m low-resolution map, we set a band integrated intensity threshold of $^{-6}$ $^{-2}$ $^{-1}$."
 For the 11.0 um high-resolution map we set a threshold of 1077 ?sr! and the 11.2 um high-resolution map has a threshold of 1076 Wm?sr-!., For the 11.0 $\mu$ m high-resolution map we set a threshold of $^{-7}$ $^{-2}$ $^{-1}$ and the 11.2 $\mu$ m high-resolution map has a threshold of $^{-6}$ $^{-2}$ $^{-1}$.
 The [6.2]/[11.2] vs [11.0]/[11.2] um ratio in NGC 7023 is presented in Figure 10.., The [6.2]/[11.2] vs [11.0]/[11.2] $\mu$ m ratio in NGC 7023 is presented in Figure \ref{fig:linear}.
" The data reveal a clear correlation, validating the use of the [11.0]/[11.2] um ratio as a PAH ionization indicator."," The data reveal a clear correlation, validating the use of the [11.0]/[11.2] $\mu$ m ratio as a PAH ionization indicator."
 The outliers in the upper left corner correlate to spectra where the thermal continuum from the source star is contaminating the linear continuum subtraction., The outliers in the upper left corner correlate to spectra where the thermal continuum from the source star is contaminating the linear continuum subtraction.
" For this reason, these points were not included in the linear fit."," For this reason, these points were not included in the linear fit."
" The linear fit has a high correlation coefficient of 0.95, from which an empirical relation can be derived: The BSS analysis identifies an independent broad component underneath the well-known 11.2 and 12.7 uum bands."," The linear fit has a high correlation coefficient of 0.95, from which an empirical relation can be derived: The BSS analysis identifies an independent broad component underneath the well-known 11.2 and 12.7 $\mu$ m bands."
" Earlier BSS studies over a wider wavelength range and the spatial distribution of Signal 3 (Figure 5)) assigns this component to emission by VSGs, proposed to be PAH clusters."," Earlier BSS studies over a wider wavelength range and the spatial distribution of Signal 3 (Figure \ref{fig:contours}) ) assigns this component to emission by VSGs, proposed to be PAH clusters."
 Early observations of the 11.2 um feature and its underlying emission support the suggestion that this broad underlying pedestal arises from a separate component (?).., Early observations of the 11.2 $\mu$ m feature and its underlying emission support the suggestion that this broad underlying pedestal arises from a separate component \citep{1985MNRAS.215..425R}.
 The PAH Database analysis provides some further insight in the character of the carrier of this broad component., The PAH Database analysis provides some further insight in the character of the carrier of this broad component.
" In this analysis, the 11-15 ym pedestal emission is due to a large number of individual components originating in a wide variety of molecular edge structures (solo's, duo's, and trio's), which together blend in an indistinct broad emission bump from 11 - 15 um. For this blend, the analysis selects relatively small species from the database."," In this analysis, the 11-15 $\mu$ m pedestal emission is due to a large number of individual components originating in a wide variety of molecular edge structures (solo's, duo's, and trio's), which together blend in an indistinct broad emission bump from 11 - 15 $\mu$ m. For this blend, the analysis selects relatively small species from the database."
" However, that is a selection effect."," However, that is a selection effect."
" Small PAHs have, by necessity, a preponderance of corner structures."," Small PAHs have, by necessity, a preponderance of corner structures."
" In contrast, calculations for large PAHs have focused (for obvious reasons) on regular structures with long straight edges and consequently strong 11.2 um bands and weak bands at longer wavelengths."," In contrast, calculations for large PAHs have focused (for obvious reasons) on regular structures with long straight edges and consequently strong 11.2 $\mu$ m bands and weak bands at longer wavelengths."
 We surmise that large irregular PAHs would equally fit the bill., We surmise that large irregular PAHs would equally fit the bill.
 The VSG component has been assigned to clusters of PAHs based upon an interpretation of the observed spatial distribution and the physical properties of clusters (???)..," The VSG component has been assigned to clusters of PAHs based upon an interpretation of the observed spatial distribution and the physical properties of clusters \citep{Olivier, Rapacioli,2006A&A...460..519R}."
" However, the spectroscopic properties of PAH clusters are presently unknown."," However, the spectroscopic properties of PAH clusters are presently unknown."
" While in general their spectra might be expected to resemble those of the constituent PAH molecules making up the cluster, we surmise that steric hindrance may affect the frequencies of the out-of-plane CH bending modes."," While in general their spectra might be expected to resemble those of the constituent PAH molecules making up the cluster, we surmise that steric hindrance may affect the frequencies of the out-of-plane CH bending modes."
 We realize that there is a hidden issue here: the spectral differences in the 11-15 um range — the broad and indistinct band in the VSG component versus the very distinct 11.2 and 12.7 modes of the PAHs - implies a more complicated evolutionary relationship between the VSGs and the PAHs than simple evaporation., We realize that there is a hidden issue here: the spectral differences in the 11-15 $\mu$ m range – the broad and indistinct band in the VSG component versus the very distinct 11.2 and 12.7 modes of the PAHs – implies a more complicated evolutionary relationship between the VSGs and the PAHs than simple evaporation.
" Applying a BSS method to observations from Spitzer's Infrared Spectrograph, Short-wavelength High-resolution mode, we uncovered 3 component signals in the PDR NGC 7023-NW."," Applying a BSS method to observations from Spitzer's Infrared Spectrograph, Short-wavelength High-resolution mode, we uncovered 3 component signals in the PDR NGC 7023-NW."
" We found that each signal is most abundant in different regions of the PDR, depending on the radiation environment."," We found that each signal is most abundant in different regions of the PDR, depending on the radiation environment."
" We identified the three component signals as PAH cations, neutral PAHs, and VSGs."," We identified the three component signals as PAH cations, neutral PAHs, and VSGs."
" As the observed spectra suggest, the neutral PAHs dominate every region of the PDR, but are most heavily"," As the observed spectra suggest, the neutral PAHs dominate every region of the PDR, but are most heavily"
a clear flare: the whole exposure (38ks for RGSI-orderl and 36ks for RGS2-ordorl) was thus kept.,a clear flare: the whole exposure (38ks for RGS1-order1 and 36ks for RGS2-order1) was thus kept.
" The overall count rates in the kkeV enerev band are 3.71+0.25&102 31 and 6.69£0.25«10.7 for the first orders of (ασ aud RGS2. respectively,"," The overall count rates in the keV energy band are $3.71 \pm 0.25 \times10^{-2}$ $^{-1}$ and $6.69 \pm 0.25 \times10^{-2}$ $^{-1}$ for the first orders of RGS1 and RGS2, respectively."
 This correspouds o net uunmiber of counts of about 1100 aud 2LOO for je first orders of RGSI ane RGS2. respectively.," This corresponds to net number of counts of about 1400 and 2400 for the first orders of RGS1 and RGS2, respectively."
 A ooupiue of the chiaunels by a factor of 2 was used. o iuprove the signal-to-noise without deeracdiug 1ο resolution.," A grouping of the channels by a factor of 3 was used, to improve the signal-to-noise without degrading the resolution."
 Note that (1) only first-order data were cousidered since the second-order data are very nolsv: (2) only RCS2 data are available at 1e waveleneth of the neon lines because of a CCD απο in RCSL., Note that (1) only first-order data were considered since the second-order data are very noisy; (2) only RGS2 data are available at the wavelength of the neon lines because of a CCD failure in RGS1.
 Both RGS aud MEG spectra are shown in Fie. l.., Both RGS and MEG spectra are shown in Fig. \ref{spec}.
" The strongest lines are ΙΔΙΑ aud ALIA Που the RGS aud ADA... AG.2QA.. AG.TA.. LAS. LA.. ,À9.2À,, and AL2A Που the ALEC."," The strongest lines are $\lambda$ and $\lambda$ for the RGS and $\lambda$, $\lambda$, $\lambda$, $\lambda$ , $\lambda$, and $\lambda$ for the MEG."
 It may be noted that lines of highlv-ionized iron are detected. which is not common for ‘normal’ O stars.," It may be noted that lines of highly-ionized iron are detected, which is not common for `normal' O stars."
 Tn Fig. 2..," In Fig. \ref{global},"
" we also compare the spectrum of Πο those of the magnetic object 0 CC (ObsID 3.0). and two ‘normal’ O stars. MMon (ObsID 5101.6217) aud 2206267 (ObsID 1888.1889). frou, the sequence of (2009)."," we also compare the spectrum of to those of the magnetic object $\theta^1$ C (ObsID 3,4), and two `normal' O stars, Mon (ObsID 5401,6247) and 206267 (ObsID 1888,1889), from the sequence of ."
. Of these four stars. MMonu suffers from the stnallest interstellar absorption (E(CD-V)20.03). while 04 O0O0nCC is more extncted with E(D-Vjy=0.29 aud the last two objects display the highest extinctions. with E(B-V)~0.5.," Of these four stars, Mon suffers from the smallest interstellar absorption (E(B-V)=0.03), while $\theta^1$ C is more extincted with E(B-V)=0.29 and the last two objects display the highest extinctions, with $\sim$ 0.5."
 This implies that differcuces between 2206267 and ccannot be due to ciffercut mterstellar absorptions. and that the hard character of 04 OOriCC is not due to a high extinction.," This implies that differences between 206267 and cannot be due to different interstellar absorptions, and that the hard character of $\theta^1$ C is not due to a high extinction."
 From Fig. 2..," From Fig. \ref{global},"
 several features are obvious., several features are obvious.
 The line luegelv dominates over the line in the spectra of 01OOn CC and1L18937.., The line largely dominates over the line in the spectra of $\theta^1$ C and.
 Though this pair of lines is the most separated one in waveleneth. lence potentially the most affected by absorption effects. this difference is not due to interstellar absorption. as aad 2206267 share similar interstellar extinctious.," Though this pair of lines is the most separated one in wavelength, hence potentially the most affected by absorption effects, this difference is not due to interstellar absorption, as and 206267 share similar interstellar extinctions."
 A siailar feature is seen for the II-to-IIe line ratio of maguesimn., A similar feature is seen for the H-to-He line ratio of magnesium.
" Reearcding silicon. the line is not clearly detected iu the “normal” O-stars. contrary toNILL. while i is of similar streneth to in anni even dominates over for 05 CC. Finally. the sulfur lines «o not appear iu the spectra of ""normal O-stars (at the scusitivity linut of the data). lxit are obvious iu 04 CC and IID115937."," Regarding silicon, the line is not clearly detected in the “normal” O-stars, contrary to, while it is of similar strength to in and even dominates over for $\theta^1$ C. Finally, the sulfur lines do not appear in the spectra of “normal” O-stars (at the sensitivity limit of the data), but are obvious in $\theta^1$ C and ."
". It should be noted however that dominates over in(δή, whereas the reverse situation is seen iu 01 CC, The reversal of the S. Si. Me and Ne IE-to- line ratios between normal stars and magnetic objects is thus clearly seen. the latter having hieher ionization."," It should be noted however that dominates over in, whereas the reverse situation is seen in $\theta^1$ C. The reversal of the S, Si, Mg and Ne H-to-He line ratios between normal stars and magnetic objects is thus clearly seen, the latter having higher ionization."
 However. it is less extreme for ttlan in the case of 01 CC. παν nevertheless participate iu the dichotoniv ciseussed by aud).," However, it is less extreme for than in the case of $\theta^1$ C. may nevertheless participate in the dichotomy discussed by and."
. The rest wavelengths of these lines were takeu romATOMDB?.. and Xspec v12.6.0.q was used or the fitting.," The rest wavelengths of these lines were taken from, and Xspec v12.6.0.q was used for the fitting."
 For IHe-like tripletsNv..SiNILL. Mex. the shifts and widths of the lines were forced to be equal for the 3 components: whereas or the ILlike doublets(Siniv..NIL.N.. Ovi). the relative lino intensities were in addition fied to their atomic data value at the )oal enissivity.," For He-like triplets, ), the shifts and widths of the lines were forced to be equal for the 3 components; whereas for the H-like doublets, ), the relative line intensities were in addition tied to their atomic data value at the peak emissivity."
 The signal-to-noise of the data prevents us from doing detailed line profile modelling: the lines were lius fitted by simple eaussiaus., The signal-to-noise of the data prevents us from doing detailed line profile modelling: the lines were thus fitted by simple gaussians.
 The fitting results are provided in Table 1 aud shown ou Figures 3. to l.., The fitting results are provided in Table \ref{linefit} and shown on Figures \ref{lineH} to \ref{lineHe}.
 For each line. the rest wavelengthfenerey. shift. width. and intensities are tabulated. together with their associated le error.," For each line, the rest wavelength/energy, shift, width, and intensities are tabulated, together with their associated $1\sigma$ error."
 Note that these errors are often asvnuuetriceal: the value shown here always is the largest value., Note that these errors are often asymmetrical: the value shown here always is the largest value.
 The last coliuun provides theintensities correctedfor aninterstellar, The last column provides theintensities correctedfor aninterstellar
regions can not satisfactorily be considered: as. separate entities. which once again has important implications for luture models that aim to examine mixing and transport in stars with multiple convection zones.,"regions can not satisfactorily be considered as separate entities, which once again has important implications for future models that aim to examine mixing and transport in stars with multiple convection zones."
 There is still a variety of questions and issues that need to be fully resolved about the complex dynamical interactions that occur within A-type stars., There is still a variety of questions and issues that need to be fully resolved about the complex dynamical interactions that occur within A-type stars.
 The mere presence of multiple, The mere presence of multiple
over the full observation period (Table 13) is 1.15 Crab. consistent with the DATSE hieh eamuna- state (59 of Lingctal.(1987. 1997))).,"over the full observation period (Table \ref{Flux_table}) ) is 1.15 Crab, consistent with the BATSE high gamma-ray state $\gamma_2$ of \citet{Ling1987,Ling1997}) )."
" The observed GDM flux ratios are Που=0.226£0.001. Rig)=0.119£0.001. and Ragy=0.01040,001. inconsistent with a siugle power luv."," The observed GBM flux ratios are $R_{50} = 0.226 \pm 0.001$, $R_{100} = 0.119 \pm 0.001$, and $R_{300} = 0.010 \pm 0.001$, inconsistent with a single power law."
 À single power hav with D—2 would vieldl flux ratios 0.158. 0.105. and 0.021.," A single power law with $\Gamma = 2$ would yield flux ratios 0.158, 0.105, and 0.021."
 Iustead. the CDM ratios sugeest a spectrum that appears significantly flatter at low cnuereies aud steeper at Ligh energies. consistent with the behavior reported earlier from. the BATSE analysis.," Instead, the GBM ratios suggest a spectrum that appears significantly flatter at low energies and steeper at high energies, consistent with the behavior reported earlier from the BATSE analysis."
 The relatively. nearby radio ealaxy Con A is a Seyfert 2 ealaxy that is the brightest ACN in hard N-ravs/low euergv goanuua rays., The relatively nearby radio galaxy Cen A is a Seyfert 2 galaxy that is the brightest AGN in hard X-rays/low energy gamma rays.
 It has powerful jets alieued at approximately 707 from the line of sieht and is seen to vary on fine scales of tens of days to vears., It has powerful jets aligned at approximately $70^\circ$ from the line of sight and is seen to vary on time scales of tens of days to years.
 It has been observed at hard A-rav cuereies byOSSE (singerctal. 1995).. audRXTE (Rothschildetal.2006).. and at enereies >1 MeV by COMPTEL (Steiule(al. 1998).," It has been observed at hard X-ray energies by \citep{Kinzer1995}, , and \citep{Rothschild2006}, and at energies $>1$ MeV by COMPTEL \citep{Steinle1998}."
. The observations below 150 keV are consistent with a lard spectrum with a power law index D.~Ls1.9., The observations below 150 keV are consistent with a hard spectrum with a power law index $\Gamma \sim 1.8-1.9$.
 The combined OSSE aud COMPTEL data are cousistent with a steepening of the spectrmm at 150 keV to P—2.3. with the spectrum then extcuding unbroken to beyond LO MeV. The CDM light curve for Cen A is shown iu Fie. L.," The combined OSSE and COMPTEL data are consistent with a steepening of the spectrum at 150 keV to $\Gamma \sim 2.3$, with the spectrum then extending unbroken to beyond 10 MeV. The GBM light curve for Cen A is shown in Fig. \ref{CenA}."
" Because Cou A is relatively fay below the equatorial plane. with a declination ó=13"". its veta angle (which ranges between à E) can be arecr than the halfanele size of the Earth as secu roni (gu& στι}."," Because Cen A is relatively far below the equatorial plane, with a declination $\delta = -43^{\circ}$, its beta angle (which ranges between $\delta \pm i$ ) can be larger than the half-angle size of the Earth as seen from $\beta_{\rm{earth}} \approx \pm 67^{\circ}$ )."
 When this happens. Cen Ais not occulted.," When this happens, Cen A is not occulted."
 This causes periodic eaps iu he light curve. with the period of the gaps equal o the precession period of the orbit.," This causes periodic gaps in the light curve, with the period of the gaps equal to the precession period of the orbit."
 The fluxes as measured by CDM aud given iu Table d. are consistent with the hard spectra. neastredl by previous iustruueuts., The fluxes as measured by GBM and given in Table \ref{Flux_table} are consistent with the hard spectrum measured by previous instruments.
 The flix ratios uecasured by CDM are Που=0.168+0.008aud, The flux ratios measured by GBM are $R_{50} =0.168 \pm 0.008$and
"precessional phase. implying variations in its density. JV... between 10!"" and 107 7.","precessional phase, implying variations in its density, $N_{\rm e}$, between $10^{10}$ and $10^{13}$ $^{-3}$."
 Thus. the physical parameters of the gaseous components imply rather dense environments emitting the Balmer lines.," Thus, the physical parameters of the gaseous components imply rather dense environments emitting the Balmer lines."
 We thank the referee for a careful reading of the paper., We thank the referee for a careful reading of the paper.
 We are very grateful to the Sciences and Technology Facilities Council (STFC) and Conieyt for the award of a STFC-Gemini studentship., We are very grateful to the Sciences and Technology Facilities Council ) and Conicyt for the award of a -Gemini studentship.
 We are specially grateful to Greg Aldering. the Nearby Supernova Factory collaboration and the University of Hawaii for their generosity in allowing us to obtain some spectra on 4433 during interstitial time within their main supernovae follow up campaign.," We are specially grateful to Greg Aldering, the Nearby Supernova Factory collaboration and the University of Hawaii for their generosity in allowing us to obtain some spectra on 433 during interstitial time within their main supernovae follow up campaign."
 K. B. thanks the Royal Society for a University Research Fellowship., K. B. thanks the Royal Society for a University Research Fellowship.
model is a better description of the data at a statistically significant level (> 3c).,model is a better description of the data at a statistically significant level $>3\sigma$ ).
" Finally, we add a third triaxial component."," Finally, we add a third triaxial component."
" In general, whether we fit three independent triaxial components or an oblate disk with two independent triaxial components, the fitting process de-weights the third component."," In general, whether we fit three independent triaxial components or an oblate disk with two independent triaxial components, the fitting process de-weights the third component."
" The resulting model is dominated by only two components, which make up >90% of the model."," The resulting model is dominated by only two components, which make up $>$ of the model."
" For the rest of the paper, we will use the two component model with one constrained to be an oblate spheroid for simplicity."," For the rest of the paper, we will use the two component model with one constrained to be an oblate spheroid for simplicity."
" For the highest mass bin, we find a low oblate spheroid fraction, 13 + 496.."," For the highest mass bin, we find a low oblate spheroid fraction, 13 $\pm$ 4."
" The best fitting model has T~0.4, near a triaxial model in the formalism of Franx (1991)."," The best fitting model has $T\simeq0.4$, near a triaxial model in the formalism of \citet{franx1991}."
". This is close to the results found in that paper, and what we expected for large, dispersion supported systems."," This is close to the results found in that paper, and what we expected for large, dispersion supported systems."
" Below 1011Mo, we can see the long, flat distribution to small axis-ratios well described by an oblate spheroid."," Below $10^{11}\ M_{\sun}$, we can see the long, flat distribution to small axis-ratios well described by an oblate spheroid."
" In the middle mass bin, the fraction of oblate spheroids is aand it grows696 to iin the lowest mass bin."," In the middle mass bin, the fraction of oblate spheroids is and it grows to in the lowest mass bin."
" Interestingly, for the middle mass"," Interestingly, for the middle mass"
44686 CCyg profile.,4686 Cyg profile.
 The presence of the 4471 line suggests an OB companion: it is quite possible that the line arises in. the close resolved companions. similarly to the case of LHI0-3209 (?)..," The presence of the 4471 line suggests an OB companion; it is quite possible that the line arises in the close resolved companions, similarly to the case of LH10-3209 \citep{walborn99}."
 However. contamination by unresolved close components can not be excluded.," However, contamination by unresolved close components can not be excluded."
 Assuming an intrinsic color of E(B—V)== --0.32 mag for early O type stars (?) and a distance modulus of mmag. the absolute magnitude of star," Assuming an intrinsic color of $E(B-V)$ –0.32 mag for early O type stars \citep{walborn02} and a distance modulus of mag, the absolute magnitude of star"
on a timescale of the stellar crossing time across the sphere of influence. ~10° vvears.,"on a timescale of the stellar crossing time across the sphere of influence, $\sim 10^3$ years."
 This research has been supported by the Australian Research Council through Discovery Project erant DP 0936336 and the National Science Foundation through grant AST-0807400., This research has been supported by the Australian Research Council through Discovery Project grant DP 0986386 and the National Science Foundation through grant AST-0807400.
 2T.Linag 2gC (FrecimanL970).. Dotluu11995:vandenTocketal.2000:Dell2001. Bolun.hupey," $\gtrsim 1$ $^{-2}$ \citep{freeman70}. \citep[e.g.][]{vdh93,mcg_opt94,mcg_abund94,mcg_gas97,edb_phot95,vdhoek00,
bell_lsb00}."
&AIcGaueh(1997) ΠΠ (Navarro.Frenk&," \citet{lsbreview} \citep{NFW96},"
in the Balmer and Hines.,in the Balmer and lines.
 This double-hump behaviour is also seen during the September 1997 outburst in the infrared (Webb 1999) and may result. from shadowing bv the spiral shocks (Stecehs priv., This double-hump behaviour is also seen during the September 1997 outburst in the infrared (Webb 1999) and may result from shadowing by the spiral shocks (Steeghs priv.
 comm.)., comm.).
 Adding the quacratie term of the ephemeris (as given by Wolf 1993) results in an orbital phase hange of ~0.006 which cannot explain the shifts observed --n the eclipses of the emission lines., Adding the quadratic term of the ephemeris (as given by Wolf 1993) results in an orbital phase change of $\sim$ 0.006 which cannot explain the shifts observed in the eclipses of the emission lines.
 By assuming a Weplerian accretion Llow we transforni Ίο velocity. coordinates of the Doppler maps to Esoe coordinates and compute model. spectra for1..A4686X... ancl X.," By assuming a Keplerian accretion flow we transform the velocity coordinates of the Doppler maps to space coordinates and compute model spectra for, and ."
.. Phese spectra. shown in 16 lower panels of Eig.," These spectra, shown in the lower panels of Fig."
 4. are caleulated at all orbital phases. including during eclipse.," 4, are calculated at all orbital phases, including during eclipse."
 Although the spectra during the eclipse were not used. to obtain the Doppler images. we can calculate the computed. spectra during eclipse.," Although the spectra during the eclipse were not used to obtain the Doppler images, we can calculate the computed spectra during eclipse."
 These spectra can be caleulated from the maps by using IP Peg's parameters. given by Alarsh Lorne (1990).," These spectra can be calculated from the maps by using IP Peg's parameters, given by Marsh Horne (1990)."
 “Phe translation from velocities to positions becomes cloubtful at low velocities which can end. outside the Roche lobe., The translation from velocities to positions becomes doubtful at low velocities which can end outside the Roche lobe.
 We arbitrarily decided to ignore any eclipse of points which mapped to locations bevond r — Ry., We arbitrarily decided to ignore any eclipse of points which mapped to locations beyond r = $_{{\rm L}1}$.
 The light curves generated from the computed ancl real data are plotted in Fig., The light curves generated from the computed and real data are plotted in Fig.
 5 with solid lines and. crosses respectively., 5 with solid lines and crosses respectively.
 The eclipses of the computed. and real data do not occur at the same time., The eclipses of the computed and real data do not occur at the same time.
 The model data do not show the shift of minimum light towards earlier orbital phases clisplaved by the data., The model data do not show the shift of minimum light towards earlier orbital phases displayed by the data.
 Shifts of the sort seen in the data require a large asvmmetry in the emission line (ux clistribution about the line of centres of the two stars., Shifts of the sort seen in the data require a large asymmetry in the emission line flux distribution about the line of centres of the two stars.
 The spiral shocks do have this asvnunetry but apparently not large enough to explain the shifts that we sce., The spiral shocks do have this asymmetry but apparently not large enough to explain the shifts that we see.
 We have no plausible explanation for this., We have no plausible explanation for this.
 We noticed in section 3.2 that the secondary star to contributedthe Balmer emission anc almost none of the emission., We noticed in section \ref{sec:doppler} that the secondary star contributed to the Balmer emission and almost none of the emission.
 In the top panels of Fig. 6..," In the top panels of Fig. \ref{res:comp},"
 we show the Doppler maps obtained from the combined: spectra for both nights but this time centered. ancl expanded: around the secondary star instead. of the centre of mass of the system., we show the Doppler maps obtained from the combined spectra for both nights but this time centered and expanded around the secondary star instead of the centre of mass of the system.
 The svnimetri¢ component of the accretion disc for hhas been subtracted and the saturation levels on the map adjusted to emphasize the companion star., The symmetric component of the accretion disc for has been subtracted and the saturation levels on the map adjusted to emphasize the companion star.
 The ratio of the masses of both components used is ¢=0.49 and the racial velocity. semi-amiplitude of the secondary A»=300kms1 , The ratio of the masses of both components used is $q = 0.49$ and the radial velocity semi-amplitude of the secondary $K_2 = 300\kmsec$ 
Figure 2a shows a V versus V.I1 colour magnitude diagram for the stars of the Irwin et al. (,Figure 2a shows a $V$ versus $V-I$ colour magnitude diagram for the stars of the Irwin et al. (
2008) sample (with rotation periods) which lie within the field of view.,2008) sample (with rotation periods) which lie within the field of view.
 As a comparison we show photometric members of NGC 2547 (from Naylor et al., As a comparison we show photometric members of NGC 2547 (from Naylor et al.
 2002) that also lie within the X-ray field of view and indicate those which have an X-ray counterpart among the 323 significant X-ray sources reported in section 2.1., 2002) that also lie within the X-ray field of view and indicate those which have an X-ray counterpart among the 323 significant X-ray sources reported in section 2.1.
 Figure 2a shows that the Irwin et al. (, Figure 2a shows that the Irwin et al. (
2008) sample is only about 50 per cent complete.,2008) sample is only about 50 per cent complete.
 A significant fraction of photometic cluster candidates at all colours have no detected rotation periods., A significant fraction of photometic cluster candidates at all colours have no detected rotation periods.
 Most of these non-periodie candidates are X-ray detected and very likely to be genuine NGC 2547 members (see Jeffries et al., Most of these non-periodic candidates are X-ray detected and very likely to be genuine NGC 2547 members (see Jeffries et al.
 2006 for a detailed discussion)., 2006 for a detailed discussion).
 Only for stars with Vo/« L5and V/>2.5 are there significant numbers of candidates without X-ray detections., Only for stars with $V-I<1.5$ and $V-I>2.8$ are there significant numbers of candidates without X-ray detections.
 For the former. these are likely to be contaminating giants among the photometrically selected members. but the latter are more likely to be genuine cluster members that become too faint for detection at very low luminosities (see below). as indeed are many of the Irwin et al. (," For the former, these are likely to be contaminating giants among the photometrically selected members, but the latter are more likely to be genuine cluster members that become too faint for detection at very low luminosities (see below), as indeed are many of the Irwin et al. ("
2008) objects at similar colours.,2008) objects at similar colours.
 Figure 2b plots X-ray activity. expressed as logLy/Liu. versus colour.," Figure 2b plots X-ray activity, expressed as $\log L_{\rm x}/L_{\rm
 bol}$, versus colour."
 There is a gradually rising mean activity level as we move from K- through to M-dwarts., There is a gradually rising mean activity level as we move from K- through to M-dwarfs.
 The upper envelope is ill defined due to three very high points that seem well separated from the rest of the distribution., The upper envelope is ill defined due to three very high points that seem well separated from the rest of the distribution.
 À time-series analysis of these three stars reveals that each was affected by a large flare during the course of theXMM-Newton observation., A time-series analysis of these three stars reveals that each was affected by a large flare during the course of the observation.
 The lower envelope is well detined for |<V.—/3 and probably represents a true floor to the level of X-ray activity in cool stars at the age of NGC 2547., The lower envelope is well defined for $1< V-I <3$ and probably represents a true floor to the level of X-ray activity in cool stars at the age of NGC 2547.
 Although our NGC 2547 sample is biased by having detected rotation periods. Irwin et al. (," Although our NGC 2547 sample is biased by having detected rotation periods, Irwin et al. ("
2008) argue that there is not a population of more slowly rotating low-mass stars and Fig.,2008) argue that there is not a population of more slowly rotating low-mass stars and Fig.
 2b also shows that photometrically selected members of NGC 2547 rotation periods. and which also have a counterpart. share a similar minimum (and median) level of X-ray activity as a function of colour.," 2b also shows that photometrically selected members of NGC 2547 rotation periods, and which also have a counterpart, share a similar minimum (and median) level of X-ray activity as a function of colour."
 Neither is this minimum level a product of the sensitivity of the X-ray observations. because almost all photometric candidates with 2.5 and which lie in the field of view have been detected (see Fig.," Neither is this minimum level a product of the sensitivity of the X-ray observations, because almost all photometric candidates with $V-I < 2.8$ and which lie in the field of view have been detected (see Fig."
" 2a),", 2a).
 The important point that emerges from these arguments is that the range of X-ray activity levels in NGC 2547 at a given colour. is quite small (less than an order of magnitude) for voung K- and M- with ο7«2.5.," The important point that emerges from these arguments is that the range of X-ray activity levels in NGC 2547 at a given colour, is quite small (less than an order of magnitude) for young K- and M-dwarfs with $V-I<2.8$."
 For cooler stars with ο>2.5. the X-ray observations become progressively incomplete and we can say very little about the range of X-ray activity here.," For cooler stars with $V-I>2.8$, the X-ray observations become progressively incomplete and we can say very little about the range of X-ray activity here."
 Figure 3 shows the dependence of activity (Ly/Li) on rotation period. considering here only those stars in NGC 2547 with rotation periods given by Irwin et al. (," Figure 3 shows the dependence of activity $L_{\rm x}/L_{\rm bol}$ ) on rotation period, considering here only those stars in NGC 2547 with rotation periods given by Irwin et al. ("
2008).,2008).
 For the purposes of later discussion we have divided the stars up according to their estimated masses: 0.55«ΑΛΛΗ.0.95AL. (roughly corresponding to stars. blue circles): 0.35<ΛΑΔΙ.0.55 (corresponding to MO-MB stars. red crosses): and AL«0.35M. (corresponding to stars cooler than M3. black open circles).," For the purposes of later discussion we have divided the stars up according to their estimated masses: $0.55<M/M_{\odot}<0.95\,M_{\odot}$ (roughly corresponding to K-stars, blue circles); $0.35<M/M_{\odot}<0.55$ (corresponding to M0-M3 stars, red crosses); and $M<0.35\,M_{\odot}$ (corresponding to stars cooler than M3, black open circles)."
 The masses were estimated from the luminosities of the stars and a MMyr isochrone from the evolutionary models of Siess. Dufour Forestini (2000).," The masses were estimated from the luminosities of the stars and a Myr isochrone from the evolutionary models of Siess, Dufour Forestini (2000)."
 The bolometric luminosities of our sample were calculated using the corrected. intrinsic V magnitudes. the bolometric corrections described in section 2.3 and an assumed distance to NGC 2547 of ppe (e.g. Mayne Naylor 2008).," The bolometric luminosities of our sample were calculated using the corrected, intrinsic $V$ magnitudes, the bolometric corrections described in section 2.3 and an assumed distance to NGC 2547 of pc (e.g. Mayne Naylor 2008)."
 The choice of the mass division at 0.55Al. is to isolate K- from the cooler M-dwarfs and at 0.35Al. to mark the approximate transition to a fully convective star (Siess et al.," The choice of the mass division at $0.55\,M_{\odot}$ is to isolate K-dwarfs from the cooler M-dwarfs and at $0.35\,M_{\odot}$ to mark the approximate transition to a fully convective star (Siess et al."
 2000)., 2000).
 lore importantly. the latter division marks the lowest mass at which our sample can be considered complete. in the sense that upper limits begin to occur in stars of lower mass.," More importantly, the latter division marks the lowest mass at which our sample can be considered complete, in the sense that X-ray upper limits begin to occur in stars of lower mass."
 The lowest uminosity stars in the sample have masses of &0.1M...," The lowest luminosity stars in the sample have masses of $\simeq 0.1\,M_{\odot}$."
 Figure 3 shows that X-rays have been detected from K- and M- with periods ranging from ddays., Figure 3 shows that X-rays have been detected from K- and M-dwarfs with periods ranging from days.
 Most of the upper imits occur for stars with short periods., Most of the upper limits occur for stars with short periods.
 Given the possibility of a correlation between X-ray activity and rotation this at first seems surprising., Given the possibility of a correlation between X-ray activity and rotation this at first seems surprising.
 The reason is that most of the lowest luminosity Cand yence lowest mass) objects in the Irwin et al. (, The reason is that most of the lowest luminosity (and hence lowest mass) objects in the Irwin et al. (
2008) catalogue have short rotation periods.,2008) catalogue have short rotation periods.
" As the the ratio Ly/L4, is close-to-constant"," As the the ratio $L_{\rm
x}/L_{\rm bol}$ is close-to-constant"
We have shown that the erating spectrum of the extreme O star. Η 93129. can be understood. using the same paracigm that explains the canonical embedded: wind shock source.Pup.. once adjustments are made for its larger wind terminal velocity and mass-loss rate.,"We have shown that the grating spectrum of the extreme O star, HD 93129A, can be understood using the same paradigm that explains the canonical embedded wind shock source, once adjustments are made for its larger wind terminal velocity and mass-loss rate."
 Specifically. the kinematies of the X-ray emitting plasma are consistent with shocks embedded in a j—Q0T. ο=3200kms wind starting at several tenthsfs... and that the attenuation signatures in the line profiles are consistent with a mass-loss rate of 4.7 to 7.0.10°-. Fepresentine. a moclest reduction compared to traditional mass-loss rates determined from measurements that ignore the ellects of clumping. and showing consistency with modeling that includes modest. clumping. in line with what is seen in LDL simulations (Dessart&Owocki2005).," Specifically, the kinematics of the X-ray emitting plasma are consistent with shocks embedded in a $\beta = 0.7$ , $\vinf
= 3200$ wind starting at several tenths, and that the attenuation signatures in the line profiles are consistent with a mass-loss rate of 4.7 to $7.0 \times 10^{-6}$, representing a modest reduction compared to traditional mass-loss rates determined from measurements that ignore the effects of clumping, and showing consistency with modeling that includes modest clumping, in line with what is seen in LDI simulations \citep{do2005}."
. This mass-loss rate reduction of a factor of three to four is consistent with that found for (Cohenetal.2010)., This mass-loss rate reduction of a factor of three to four is consistent with that found for \citep{Cohen2010}.
. We have also demonstrated: for he first time that. modeling wind. absorption of N-ravs or both line profiles and. for the broadband. spectral energev distribution leads to consistent results when a »hysically realistic model of the broadband wind attenuation (Leuteneggerοἱal.2010). is used., We have also demonstrated for the first time that modeling wind absorption of X-rays for both line profiles and for the broadband spectral energy distribution leads to consistent results when a physically realistic model of the broadband wind attenuation \citep{Leutenegger2010} is used.
 The global spectral modeling indicates. that the dominant thermal emission component has quite a modest emperature. of roughly 0.0 keV. as. predicted. by EWS mocels.," The global spectral modeling indicates that the dominant thermal emission component has quite a modest temperature, of roughly 0.6 keV, as predicted by EWS models."
 Ehe observed. overall hardness of the spectrum is attributable to wind attenuation. rather than high plasma emperatures.," The observed overall hardness of the spectrum is attributable to wind attenuation, rather than high plasma temperatures."
 Phe low dominant plasma temperature is also manifest in the low ratio. which is consistent with the value found in the erating spectrum ofPup.," The low dominant plasma temperature is also manifest in the low ratio, which is consistent with the value found in the grating spectrum of."
. There is likely a small amount of hard. X-ray emission from colliding wind binary interaction between components Aa and Ab. many tens of AU from either stars photosphere.," There is likely a small amount of hard X-ray emission from colliding wind binary interaction between components Aa and Ab, many tens of AU from either star's photosphere."
 Because of the large separation of the components. this X-ray. emission makes a small contribution to the overall X-ray spectral properties. representing less than of the systems X-ray luminosity.," Because of the large separation of the components, this X-ray emission makes a small contribution to the overall X-ray spectral properties, representing less than of the system's X-ray luminosity."
 The helium-like f/; ratios also provide evidence that the bulk of the X-rays arise in embedded wind shocks., The helium-like $f/i$ ratios also provide evidence that the bulk of the X-rays arise in embedded wind shocks.
 As HD 931294 is the earliest O star known. and has one of the strongest winds of any O star. the work presented here strongly suggests that the embedded: wind shock scenario. as described. by numerical. simulations of the line-driving instability. is widely applicable to O stars. even those with extremely strong winds.," As HD 93129A is the earliest O star known, and has one of the strongest winds of any O star, the work presented here strongly suggests that the embedded wind shock scenario, as described by numerical simulations of the line-driving instability, is widely applicable to O stars, even those with extremely strong winds."
 nd that X-ray line profile analysis. especially in conjunction with broadband spectral modeling. provides a good means of making a clumping-independent niass-loss rate determination for O stars with dense winds.," And that X-ray line profile analysis, especially in conjunction with broadband spectral modeling, provides a good means of making a clumping-independent mass-loss rate determination for O stars with dense winds."
 Support [for this work was provided by the National Acronautics and Space Administration through award numbers ART-s002N and €O0-11002D to swarthmoreo College., Support for this work was provided by the National Aeronautics and Space Administration through award numbers AR7-8002X and GO0-11002B to Swarthmore College.
 EEW was supported by a Lotte Lazarsfelcdl Bailvn Summer Research Fellowship ancl JPAL was supported. by a Surdna Summer Research Fellowship. both from the Provost's Ollice at Swarthmore College.," EEW was supported by a Lotte Lazarsfeld Bailyn Summer Research Fellowship and JPM was supported by a Surdna Summer Research Fellowship, both from the Provost's Office at Swarthmore College."
 ALAL is supported by an appointment to the NASA Postdoctoral Program at Cioddard Space Flight Center. administered by Oak Ridge Associated Universities through a contract with NASA.," MAL is supported by an appointment to the NASA Postdoctoral Program at Goddard Space Flight Center, administered by Oak Ridge Associated Universities through a contract with NASA."
 JOS and SPO acknowledge support [rom NASA award AVP NNXILACAROCG to the University of Delaware., JOS and SPO acknowledge support from NASA award ATP NNX11AC40G to the University of Delaware.
 The authors thank Vérronique Petit for her careful reading of themanuscript ancl several useful suggestions., The authors thank Vérronique Petit for her careful reading of themanuscript and several useful suggestions.
Brillouin (WEB) approximation in §3. summarizing the panode and e-mocle properties during the helium core flash.,"Brillouin (WKB) approximation in 3, summarizing the p-mode and g-mode properties during the helium core flash."
 We close in 84 bv discussing how stars in (his phase of evolution can be differentiated [rom (the more numerous populations of RGB and clamp stars., We close in 4 by discussing how stars in this phase of evolution can be differentiated from the more numerous populations of RGB and clump stars.
 Evolved stars with AM<2.0.7. undergo a He core flash that ends their RGB evolution. ad leads (hem to the red clump where they undergo core Le burningfor zz70 Mrs.," Evolved stars with $M\lesssim 2.0 M_\odot$ undergo a He core flash that ends their RGB evolution, and leads them to the red clump where they undergo core He burningfor $\approx 70$ Myrs."
 The flash starts as an off-center thermally unstable ignition of helium burning al a mass coordinate about hallway through the A.z0.43. Ile core 2011).. nearly independent of metallicity aud 1.," The flash starts as an off-center thermally unstable ignition of helium burning at a mass coordinate about half-way through the $M_c\approx 0.48M_\odot$ He core \citep{dominguez99,salaris02,serenelli05,serenellif05,mocak09, paxton11}, nearly independent of metallicity and $M$."
 Using MESA version 3709. we have calculated (his evolution for a range of M. all with and Z=0.02. treating convection with the Schwarzshild criterion (i.e. no semiconvection. thermohaline mixing or convective overshoot) aud a mixing leneth parameter of ay=2.0.," Using $\MESA$ version 3709, we have calculated this evolution for a range of $M$, all with $X=0.7$ and $Z=0.02$, treating convection with the Schwarzshild criterion (i.e. no semiconvection, thermohaline mixing or convective overshoot) and a mixing length parameter of $\alpha_{\rm MLT}=2.0$."
 We had no mass loss. diffusion. or rotation and we used the MESA “basic” nuclear network with rates from NACHRE.," We had no mass loss, diffusion, or rotation and we used the $\MESA$ “basic"" \citep{paxton11} nuclear network with rates from NACRE."
 The opacity was trom OPAL (Iglesias&Rogers 1996).. with the low temperature opacities taken from Fereusonοἱal.(2005) will metal ratios given by Grevesse&Sauval(1998).. as described in Paxtonetal. (2011)..," The opacity was from OPAL \citep{iglesias96}, , with the low temperature opacities taken from \citet{ferguson05} with metal ratios given by \citet{grevesse98}, as described in \citet{paxton11}. ."
 We used ihe OPAL equation-ol-state (Rogers&Navlonov2002).. and IIELM extensions BSweslv2000) where ποσο.," We used the OPAL equation-of-state \citep{rogers02}, and HELM extensions \citep{timmes00} where needed."
 The evolutionary tracks in Figure 1. fora M=1LOAL. and Af=1.511. star are lor the time before. during. ancl alter the core flash phase.," The evolutionary tracks in Figure \ref{fig:hrevol} for a $M=1.0M_\odot$ and $M=1.8M_\odot$ star are for the time before, during, and after the core flash phase."
 The He core expansion from the initiating flash leads to an adiabatic temperature drop in the overlving IE burning shell: quenching the burning., The He core expansion from the initiating flash leads to an adiabatic temperature drop in the overlying H burning shell; quenching the burning.
 This loss of an energy source lor the red giant envelope triggers a Ixelvin-IHelmnholtz (XID) contraction from the tip of the RGB on a rapid timescale (Thomas1967).. reaching Le100—200L. in 10 vears (the open blue circles).," This loss of an energy source for the red giant envelope triggers a Kelvin-Helmholtz (KH) contraction from the tip of the RGB on a rapid timescale \citep{thomas67}, reaching $L \approx 100-200 L_\odot$ in $10^4$ years (the open blue circles)."
 Thomas(1967) showed that the Ile core flash phase is governed by a series of shell subllashes Chat increase the entropy of the convective regions., \citet{thomas67} showed that the He core flash phase is governed by a series of shell subflashes that increase the entropy of the convective regions.
 This is seen in Figure 2.. where five successive subllashes are seen to move inwards in mass coordinates.," This is seen in Figure \ref{fig:structure}, where five successive subflashes are seen to move inwards in mass coordinates."
 However. convective He burning is present for onlv about of the time.," However, convective He burning is present for only about of the time."
 The /;zz2 Myr timescale is set bv the need for thermal diffision to act inwards between sublIashes. ending when the thermal wave reaches the core (Thomas1967:Serenelli&Weiss 2005).. heating it at nearly constant pressure (Paxtonetal.2011) to thecondition needed for the convective Ie burning core phaseof the red chunp.," The $t_f\approx 2$ Myr timescale is set by the need for thermal diffusion to act inwards between subflashes, ending when the thermal wave reaches the core \citep{thomas67, serenelli05}, , heating it at nearly constant pressure \citep{paxton11} to thecondition needed for the convective He burning core phaseof the red clump."
 We end Figure 2 at that time., We end Figure \ref{fig:structure} at that time.
 Even though the He, Even though the He
across the core of the quasar in the VLBA data.,across the core of the quasar in the VLBA data.
 The size scale and low velocity dispersion of the gas causing the main absorption feature suggest that it. belongs to a siugle cold cloud. or cloud complex., The size scale and low velocity dispersion of the gas causing the main absorption feature suggest that it belongs to a single cold cloud or cloud complex.
 More sensitive VLBA measurements would be necessary to determine if the weak secondary absorption is also present over this size scale. or is part of a sinaller scale feature.," More sensitive VLBA measurements would be necessary to determine if the weak secondary absorption is also present over this size scale, or is part of a smaller scale feature."
 The coustancy of the primary line characteristics over the slightly resolved core in the VLBA cata suggests that it is completely covered {ως=1) by the absorbing gas in that component. aud we assume this is true for the secouc weaker component as well.," The constancy of the primary line characteristics over the slightly resolved core in the VLBA data suggests that it is completely covered $f_\mathrm{core} = 1$ ) by the absorbing gas in that component, and we assume this is true for the second weaker component as well."
 We therefore adopt the view tliat he absorbing gas eutirely covers the core but not the exteucded weak lobes. aud that the covering actor ofthe gas is fzz0.93.," We therefore adopt the view that the absorbing gas entirely covers the core but not the extended weak lobes, and that the covering factor of the gas is $f \approx 0.98$."
 The extreme narrowness of the absorption features (as measured in the Arecibo data) places inm upper limits ou the amount of thermal broadenius iu the lines. auc consequently oui any turbulence or bulk motious of the absorbing eas.," The extreme narrowness of the absorption features (as measured in the Arecibo data) places firm upper limits on the amount of thermal broadening in the lines, and consequently on any turbulence or bulk motions of the absorbing gas."
 As a result. au upper limit to the kinetic temperature or each cloud can be found directly [rom the widths of the lines. without having to rely on a comparison between the DLA aud 2lem liue characteristics to calculate a harmonic mean spiu eniperature for the entire ensemble of clouds which lie on the sightliue.," As a result, an upper limit to the kinetic temperature for each cloud can be found directly from the widths of the lines, without having to rely on a comparison between the DLA and 21cm line characteristics to calculate a harmonic mean spin temperature for the entire ensemble of clouds which lie on the sightline."
 The kinetic temperature ol the gas is related to the width of the absorption line by: where Ae is the FWHM velocity measured in kins 1., The kinetic temperature of the gas is related to the width of the absorption line by: where $\Delta v$ is the FWHM velocity measured in km $^{-1}$.
 A simultaneous two component fit to the Arecibo spectrum left large residuals and. had a reduced 4?=2.3L, A simultaneous two component fit to the Arecibo spectrum left large residuals and had a reduced $\chi^2 = 2.34$.
 By adding a third component to the fit. the reduced 4?=1.02. aud the residuals were all within the noise.," By adding a third component to the fit, the reduced $\chi^2 = 1.02$, and the residuals were all within the noise."
 The first or main compouent has a velocity width at half maximum (FWHM) of Ae=3.68720.019 kins IF and an optical depth of 7=0.2162+0.0010.," The first or main component has a velocity width at half maximum (FWHM) of $\Delta v = 3.687 \pm 0.019$ km $^{-1}$, and an optical depth of $\tau = 0.2462 \pm 0.0010$."
 Π lies at a heliocentric frequency «C 1301.6196 MHz. corresponding to z=0.00123.," It lies at a heliocentric frequency of 1301.6496 MHz, corresponding to $z =
0.09123$."
 The second. weaker line is separated [rom the first by AV=7.002:0.01 kan 1.," The second, weaker line is separated from the first by $\Delta V_\mathrm{offset} = 7.69 \pm 0.04$ km $^{-1}$ ."
 has à FWHM velocity of Ae—2482041 km | and an optical depth 7=0.0253+0.0010., It has a FWHM velocity of $\Delta v = 2.18 \pm 0.11$ km $^{-1}$ and an optical depth $\tau = 0.0253 \pm 0.0010$.
 The third component has an optical depth 7=0.0063+0.0008 aud FWHM velocity Xe=15.2£1.L kms |.," The third component has an optical depth $\tau = 0.0063 \pm 0.0008$ and FWHM velocity $\Delta v = 15.2
\pm 1.4$ km $^{-1}$."
" The absorption is shifted by NV,=—1.5940.66 kan | with respect to the main narrow absorption feature."," The absorption is shifted by $\Delta
V_\mathrm{offset} = -1.59 \pm 0.66$ km $^{-1}$ with respect to the main narrow absorption feature."
 Attempts to force the third component to lie at the position of either of the other two components resultecl in much poorer fits., Attempts to force the third component to lie at the position of either of the other two components resulted in much poorer fits.
 Figure | shows the Arecibo spectrum after the fits to the first two components have been removed., Figure 4 shows the Arecibo spectrum after the fits to the first two components have been removed.
 The fit to the third component has been marked. and the residuals alter removing all tliree gaussian fits are shown.," The fit to the third component has been marked, and the residuals after removing all three gaussian fits are shown."
 Parameters for each of the absorption compouents are suminarized in Table 1., Parameters for each of the absorption components are summarized in Table 1.
 Usiug equation (3). the kinetic temperature is τν«p29743 Is [or the main component. and Tix«103160 Is for the secoudary line.," Using equation (3), the kinetic temperature is $T_{\mathrm k} \leq 297 \pm
3$ K for the main component and $T_{\mathrm k} \leq 103 \pm 10$ K for the secondary line."
 Both ofthese temperatures would fall within the scatter, Both ofthese temperatures would fall within the scatter
We [ind a solution of the homogeneous form of eq. B2::,We find a solution of the homogeneous form of eq. \ref{eqg}:
 Now we seek a solution of B2 of the form ἐν)=Cipyectp): whose solution is:," Now we seek a solution of \ref{eqg} of the form $g_u(p)=\tilde{C}(p)
w_C(p)$: whose solution is:"
JI890. used. morphological criteria to arrive at their original membership classification (e.g. low-Iuminosity spiral galaxies are likely to be background objects).,FS90 used morphological criteria to arrive at their original membership classification (e.g. low-luminosity spiral galaxies are likely to be background objects).
 Since we have obtained recession velocities for 103 of their 162 55044 eroup galaxies we are able to re-evaluate their. original eroup membership., Since we have obtained recession velocities for 103 of their 162 5044 group galaxies we are able to re-evaluate their original group membership.
 E90 organised their observations into three cilferent membership classes designated 1. (definite members). 2 (likely members) and 3 (possible members).," FS90 organised their observations into three different membership classes designated 1 (definite members), 2 (likely members) and 3 (possible members)."
 Our spectroscopic sample includes 60. 14 and 29 class 1. 2 and 3 galaxies respectively.," Our spectroscopic sample includes 60, 14 and 29 class 1, 2 and 3 galaxies respectively."
 Confirmed 55044 eroup members and their P890 classifications are shown in Table l1., Confirmed 5044 group members and their FS90 classifications are shown in Table \ref{group_gals}.
 Non-members. their. velocities anc E890. classifications are given in Table 2..," Non-members, their velocities and FS90 classifications are given in Table \ref{non_group}."
 Of class 1 galaxies we find 55 of the 60 to be genuine group members. 92 percent. indicating hat class 1 galaxies in the FS9O catalogue are generally reliable group member galaxies.," Of class 1 galaxies we find 55 of the 60 to be genuine group members, $\sim$ 92 percent, indicating that class 1 galaxies in the FS90 catalogue are generally reliable group member galaxies."
 Phe fractions of genuine eroup member class 2 and 3 galaxies fall. as one might expect sed on their increasing uncertainty in group membership. and we confirm 7 class 2 and 9 class 3 galaxies as bona ide group members (50 and ο] percent respectively).," The fractions of genuine group member class 2 and 3 galaxies fall, as one might expect based on their increasing uncertainty in group membership, and we confirm 7 class 2 and 9 class 3 galaxies as bona fide group members (50 and $\sim$ 31 percent respectively)."
 We note that the 55044. group. is the most. distant considered by S90. and so it is perhaps not surprising that he classification of galaxies is somewhat unreliable given he difficulty in morphologically classifving faint ealaxics.," We note that the 5044 group is the most distant considered by FS90, and so it is perhaps not surprising that the classification of galaxies is somewhat unreliable given the difficulty in morphologically classifying faint galaxies."
 Fig., Fig.
 6 shows a comparison of the B-banc luminosity. unctions (LE) between our confirmed. group sample ancl he complete FS90 sample., \ref{b_lumf} shows a comparison of the B-band luminosity functions (LF) between our confirmed group sample and the complete FS90 sample.
 The largest cillerence between he two comes at faint magnitudes. where our spectral observations become significantly incomplete. anc E890 ikelv includes. a significant number of [aint class 3 membership galaxies. of which our observations suggest onlv one-third are actually group members.," The largest difference between the two comes at faint magnitudes where our spectral observations become significantly incomplete and FS90 likely includes a significant number of faint class 3 membership galaxies, of which our observations suggest only one-third are actually group members."
 We adopt a Monte. Carlo approach to correct. the LEs shown in Fig., We adopt a Monte Carlo approach to correct the LFs shown in Fig.
 6 for our own incompleteness and contamination present in the [890 sample., \ref{b_lumf} for our own incompleteness and contamination present in the FS90 sample.
 Confirmed eroup member galaxies are always included in the analysis. however there are 59 dwarf galaxies in the P590 sample with no recession velocity information.," Confirmed group member galaxies are always included in the analysis, however there are 59 dwarf galaxies in the FS90 sample with no recession velocity information."
 We therefore sample these remaining FS9O dwarf galaxies based on the fractions of class 1. 2 and 3 included in our confirmed group sample (i.e. 92. 50 and 31 percent of cach respective membership class ave included in the Monte Carlo group cach iteration).," We therefore sample these remaining FS90 dwarf galaxies based on the fractions of class 1, 2 and 3 included in our confirmed group sample (i.e. 92, 50 and 31 percent of each respective membership class are included in the Monte Carlo group each iteration)."
 This provides a more accurate estimate of the luminosity function and allows us to examine the effect of contamination on the 178590 sample., This provides a more accurate estimate of the luminosity function and allows us to examine the effect of contamination on the FS90 sample.
 Fig., Fig.
 6. shows that our spectroscopic sample is in excellent; agreement with the P590. sample for Ale< 15. where the 900 sample is greater than 90. percent complete. and implies a [at luminosity. function.," \ref{b_lumf} shows that our spectroscopic sample is in excellent agreement with the FS90 sample for $_B<-15$ , where the FS90 sample is greater than 90 percent complete, and implies a flat luminosity function."
 1n the region 15 «Mg«14 our spectroscopic sample begins to deviate from. the E590. sample. however our. Monte Carlo corrections show that the majority of this is due to the inclusion of outliers in the P590 sample. rather than incompleteness in our data.," In the region $-15<$ $_B<-14$ our spectroscopic sample begins to deviate from the FS90 sample, however our Monte Carlo corrections show that the majority of this is due to the inclusion of outliers in the FS90 sample, rather than incompleteness in our data."
 Below Mg|14 the LE shows tentative evidence for an upturn. however our sample is significantly incomplete in this magnitude range so we abstain from drawing any further conclusions.," Below $_B>-14$ the LF shows tentative evidence for an upturn, however our sample is significantly incomplete in this magnitude range so we abstain from drawing any further conclusions."
 A vital step in establishing the role that. groups play in the evolution of galaxies is establishing their place between clusters and the field., A vital step in establishing the role that groups play in the evolution of galaxies is establishing their place between clusters and the field.
 Are groups simply low-mass clusters. or do their member galaxies exhibit unique signs of formation and evolution?," Are groups simply low-mass clusters, or do their member galaxies exhibit unique signs of formation and evolution?"
 In order to assess this we first need to establish the relation of the 55044. group to other groups and. clusters., In order to assess this we first need to establish the relation of the 5044 group to other groups and clusters.
 The dynamical properties of the 55044 group are summzrised in Table 3.. which shows the mean velocity. velocity dispersion. zooradius’. dynamical mass and crossing time we derive using the formulae in Appencix A..," The dynamical properties of the 5044 group are summarised in Table \ref{lit_param}, which shows the mean velocity, velocity dispersion, $_{500}$, dynamical mass and crossing time we derive using the formulae in Appendix \ref{Dynamical Formulae}."
 For comparison. we also show similar measurements from he literature.," For comparison, we also show similar measurements from the literature."
 While we have more than tripled the number of group members included relative to previous studies. our results or the mean velocity. velocity dispersion. Rey radius. massand crossing time are all consistent within errors.," While we have more than tripled the number of group members included relative to previous studies, our results for the mean velocity, velocity dispersion, $_{500}$ radius, massand crossing time are all consistent within errors."
 We find he group to be massive. nearly 107 MM. in total mass. which combined with the short crossing time leads us to he same conclusions of BOG that the 55044 group is a classically massive. dvnamically mature group.," We find the group to be massive, nearly $10^{14}$ $_{\odot}$ in total mass, which combined with the short crossing time leads us to the same conclusions of B06 that the 5044 group is a classically massive, dynamically mature group."
 Addressing the dynamical state of the 55044 group is key in our goal of establishing its evolutionary history., Addressing the dynamical state of the 5044 group is key in our goal of establishing its evolutionary history.
 1n particular. the brightest group galaxy 55044) is known to have a peculiar. velocity of ~150kms twith respect to the mean group velocity (Cellone Buzzoni 2005: BOG). which could. be indicative of non-equilibrium due to recent mergers.," In particular, the brightest group galaxy 5044) is known to have a peculiar velocity of $\sim$ with respect to the mean group velocity (Cellone Buzzoni 2005; B06), which could be indicative of non-equilibrium due to recent mergers."
 The peculiar velocity found for 55044. scaled by the group's velocity dispersion. is 0.45... and within the range twpically Found. for relaxed. groups and. clusters (e.g.," The peculiar velocity found for 5044, scaled by the group's velocity dispersion, is $\sigma_\nu$ , and within the range typically found for relaxed groups and clusters (e.g."
 , 
our results for daytime is consistent with that obtained by Erasmus&VanStaden(2002).,our results for daytime is consistent with that obtained by \citet{Erasmus02}.
. The differences in the values for D1 and D2 and those obtained by Erasmus&VanStaden(2002) are between 3 and 28 per cent., The differences in the values for D1 and D2 and those obtained by \citet{Erasmus02} are between 3 and 28 per cent.
" As a reference, in the case of the PWV, the results for the TMT preliminary studies carried out by Erasmus&VanStaden(2002) and those obtained from the in situ site testing group are within 30 per cent for all the sites, see Otárolaetal.(2010)."," As a reference, in the case of the PWV, the results for the TMT preliminary studies carried out by \citet{Erasmus02} and those obtained from the in situ site testing group are within 30 per cent for all the sites, see \citet{Otarola10}."
. Our results for D1 D2 are within that range., Our results for D1 D2 are within that range.
" Futhermore, the differences might be influenced by the distinct data coverage and by the definitions of clear time used by Erasmus&VanStaden (2002)."," Futhermore, the differences might be influenced by the distinct data coverage and by the definitions of clear time used by \citet{Erasmus02}."
. The results presented in this paper are based on direct solar radiation measurements., The results presented in this paper are based on direct solar radiation measurements.
 The variation trend in the centre of the clear peak shown in Fig., The variation trend in the centre of the clear peak shown in Fig.
 8 can be interpreted in terms of seasonal variations of the atmospheric transmission: during the summer months there is more atmospheric absorption than in the rest of the year., \ref{centers} can be interpreted in terms of seasonal variations of the atmospheric transmission: during the summer months there is more atmospheric absorption than in the rest of the year.
" This is consistent with the seasonal variation of the Precipitable Water Vapor (PWV) at 210 GHz reported by Hiriartetal.(1997),, Hiriartetal. (2003),, Otárolaetal.(2009, 2010)."," This is consistent with the seasonal variation of the Precipitable Water Vapor (PWV) at 210 GHz reported by \citet{Hiriart97}, \citet{Hiriart03}, \citet{Otarola09,Otarola10}."
. The seasonal variation in the PWV is shown in Figs., The seasonal variation in the PWV is shown in Figs.
" 9, 10 and 11 of Otárolaetal.(2009),, where the maximum PWV values occur during the Summer."," 9, 10 and 11 of \citet{Otarola09}, where the maximum PWV values occur during the Summer."
" The larger value of the centres of the clear peak for July 2006 and August 2008 relative to the same months of the other years, shown in Fig. 8,,"," The larger value of the centres of the clear peak for July 2006 and August 2008 relative to the same months of the other years, shown in Fig. \ref{centers},"
 suggest that the atmosphere was more transparent., suggest that the atmosphere was more transparent.
 We analysed the aerosol optical thickness reported by Araiza&Cruz-González(2011) (c.f.," We analysed the aerosol optical thickness reported by \citet{Araiza11}
 (c.f."
 Table 4)., Table 4).
 The larger value of the centre for July 2006 is, The larger value of the centre for July 2006 is
The variety of possible formation scenarios of the [Neu]] line and the large range of stellar and disk properties. including the presence of jets and outflows in some of the systems. calls for a deeper investigation requiring a much larger sample.,"The variety of possible formation scenarios of the ] line and the large range of stellar and disk properties, including the presence of jets and outflows in some of the systems, calls for a deeper investigation requiring a much larger sample."
 To this end. we have started a comprehensive study of [Neu] emission from classical T Tauri stars (CTTS) in particular in the context of X-ray emission and measured mass aceretion rates.," To this end, we have started a comprehensive study of ] emission from classical T Tauri stars (CTTS) in particular in the context of X-ray emission and measured mass accretion rates."
 Our goal has been to revisit recent work that was based nostly on small samples or observations of individual objects. and to analyze and study the combined sample of objects in à systematic way.," Our goal has been to revisit recent work that was based mostly on small samples or observations of individual objects, and to analyze and study the combined sample of objects in a systematic way."
 We have re-analyzed many of the previously reported n]|] data (mostly. from Spitzer) coherently. but we have also added crucial new observations from our dedicated observing programs.," We have re-analyzed many of the previously reported ] data (mostly from ) coherently, but we have also added crucial new observations from our dedicated observing programs."
 The present work also for the first time presents a uniform. archival study of all available X-ray data from the and observatories for these sources. complemented with data from new observing programs.," The present work also for the first time presents a uniform, archival study of all available X-ray data from the and observatories for these sources, complemented with data from new observing programs."
 The increased sensitivity and spectral resolving power of these X-ray devices permit à much better characterization of the stellar X-ray radiation than hitherto possible with data from. e.g. ROSAT as used in earlier [Να] studies (Lahuisetal...2007;Pascucci2007).," The increased sensitivity and spectral resolving power of these X-ray devices permit a much better characterization of the stellar X-ray radiation than hitherto possible with data from, e.g., ROSAT as used in earlier ] studies \citep{lahuis07, pascucci07}."
. We will further use ancillary data collected from the published literature. such as mass accretion rates or mass outflow rates.," We will further use ancillary data collected from the published literature, such as mass accretion rates or mass outflow rates."
 We also refer to forthcoming work by Baldovin-Saavedra(2010) in which a sample of gas lines (of H>. u]]. [Nem]. [Feu]]. u]D) in the IRS spectal range is discussed but for a more confined sample of pre-main sequence stars in Taurus; few significant correlations are reported there. the most likely one again supporting an origin of [Neu]] emission in outflows.," We also refer to forthcoming work by \citet{baldovin10} in which a sample of gas lines (of $_2$, ], ], ], ]) in the IRS spectal range is discussed but for a more confined sample of pre-main sequence stars in Taurus; few significant correlations are reported there, the most likely one again supporting an origin of ] emission in outflows."
 Initial results of our work are described in Güdeletal. (2009a)., Initial results of our work are described in \citet{guedel09a}.
". In short. indications of a correlation of Linewy with both Lx and Af, were found. but strong scatter dominates these correlations."," In short, indications of a correlation of $L_{\rm [Ne~II]}$ with both $L_{\rm X}$ and $\dot{M}_{\rm acc}$ were found, but strong scatter dominates these correlations."
 However. it was also found that show consistently higher Linejj. and that Linejj seems to correlate with wind or outflow properties.," However, it was also found that show consistently higher $L_{\rm [Ne~II]}$, and that $L_{\rm [Ne~II]}$ seems to correlate with wind or outflow properties."
 The purpose of the present paper is to present our entire data set and additional correlation studies. and to coherently discuss these results.," The purpose of the present paper is to present our entire data set and additional correlation studies, and to coherently discuss these results."
 Given the presently favored models for [Neu]] line emission. we selected targets for our study that have well-observed disks and may also be engines of jets an outflows. but we did not include Class I objects in which à number of additional circumstellar regions may be relevant for [Neu]] emission. such as shocks on disks produced by material accreting from the envelope. the irradiated envelopes themselves. shocks between the jets or outflows and the envelopes. ete.," Given the presently favored models for ] line emission, we selected targets for our study that have well-observed disks and may also be engines of jets an outflows, but we did not include Class I objects in which a number of additional circumstellar regions may be relevant for ] emission, such as shocks on disks produced by material accreting from the envelope, the irradiated envelopes themselves, shocks between the jets or outflows and the envelopes, etc."
 Further. strong extinction and photoelectric absorption make Class I objects difficult for study.," Further, strong extinction and photoelectric absorption make Class I objects difficult for study."
 Sufhciently strong extinction may suppress some infrared emission from the regions close to the star selectively., Sufficiently strong extinction may suppress some infrared emission from the regions close to the star selectively.
 Also. only a moderate fraction of Class | sources is accessible by modern X-ray satellites. and the detected X-ray emission is. due to photoelectric absorption. relatively hard (photon energies typically >2 keV).," Also, only a moderate fraction of Class I sources is accessible by modern X-ray satellites, and the detected X-ray emission is, due to photoelectric absorption, relatively hard (photon energies typically $>$ 2 keV)."
 The bulk part of the X-ray emission therefore remains undetected. and an unbiased reconstruction of the underlying (intrinsic) X-ray emission relies on assumptions.," The bulk part of the X-ray emission therefore remains undetected, and an unbiased reconstruction of the underlying (intrinsic) X-ray emission relies on assumptions."
 For a study predominantly based on [Neu]] detections of Class [ sources. see Flaccomioetal.(2009).," For a study predominantly based on ] detections of Class I sources, see \citet{flaccomio09}."
 Our targets were therefore required to be essentially Class IL objects or (accreting) CTTS., Our targets were therefore required to be essentially Class II objects or (accreting) CTTS.
 More massive Herbig stars were not considered given their considerable UV radiation fields and possibly very different X-ray source properties (Telleschietal..2007a)., More massive Herbig stars were not considered given their considerable UV radiation fields and possibly very different X-ray source properties \citep{telleschi07a}.
. On the other hand. CTTS ejecting jets were intentionally. included. because jets may play a crucial role in strongly aecreting CTTS. anc they may be directly linked to the accretion process. itself.," On the other hand, CTTS ejecting jets were intentionally included because jets may play a crucial role in strongly accreting CTTS, and they may be directly linked to the accretion process itself."
 Including such objects will therefore allow us to investigate to what extent jets matter for the observed [| emission. and perhaps to identify a subset of objects showing a baseline u]]| flux unaffected by Jets.," Including such objects will therefore allow us to investigate to what extent jets matter for the observed ] emission, and perhaps to identify a subset of objects showing a baseline ] flux unaffected by jets."
 To study this latter possibility further. a few targets revealing signatures of transition. disks. Le.. disks with inner holes. have been included.," To study this latter possibility further, a few targets revealing signatures of transition disks, i.e., disks with inner holes, have been included."
 Transition disks may also be important to discriminate between X-ray and EUV-related [| emission models as the potentially low gas content in the inner disk may make this region transparent to direct EUV radiation (see. e.g.. numerical models by Alexanderetal. 2006)).," Transition disks may also be important to discriminate between X-ray and EUV-related ] emission models as the potentially low gas content in the inner disk may make this region transparent to direct EUV radiation (see, e.g., numerical models by \citealt{alexander06}) )."
 The data selection is primarily driven by the availability of [Nemn]] observations (detections or upper limits)., The data selection is primarily driven by the availability of ] observations (detections or upper limits).
 Our sample is therefore mostly drawn from observations available in the IRS data archive., Our sample is therefore mostly drawn from observations available in the IRS data archive.
 The largest part of our target list originates from the Lahuisetal.(2007) survey (based on theSpitzer legacy program: Evansetal. 2003))., The largest part of our target list originates from the \citet{lahuis07} survey (based on the legacy program; \citealt{evans03}) ).
 This survey focuses on the NGC 1333. Chamaeleon. Lupus. Rho Oph. and Serpens star forming regions. all with characteristic ages of a few Myr.," This survey focuses on the NGC 1333, Chamaeleon, Lupus, Rho Oph, and Serpens star forming regions, all with characteristic ages of a few Myr."
 Compared to the preliminary presentation m Giideletal.(2009a).. we have reduced these data again using a new. improved software version with more careful. background subtraction and treatment of potential blends. resulting in many additional detections not used in the initial report.," Compared to the preliminary presentation in \citet{guedel09a}, we have reduced these data again using a new, improved software version with more careful background subtraction and treatment of potential blends, resulting in many additional detections not used in the initial report."
 To this sample. we added several targets from ourobserver. programs (PI. J. Carr).," To this sample, we added several targets from our programs (PI J. Carr)."
 This sample in particular. includes objects from the Taurus star forming region. and some targets from Chamaeleon and Rho Oph.," This sample in particular includes objects from the Taurus star forming region, and some targets from Chamaeleon and Rho Oph."
 Further u]| fluxes or upper limits thereof were adopted from the published literature. in particular for RX J1111.7-7620. PZ99 JIOI411. RX J1842.9-3542. and RX J1852.3-3700 (Pascuceietal..2007).. observed as part of the legacy program (Meyeretal.2004).. DP Tau observed with the spectrograph at Gemini. North (Herezeg 2007). and T Tau N and 5 observed with the spectrograph at the Very Large Telescope (VLT) (vanBoekeletal.. 2009).," Further ] fluxes or upper limits thereof were adopted from the published literature, in particular for RX J1111.7-7620, PZ99 J161411, RX J1842.9-3542, and RX J1852.3-3700 \citep{pascucci07}, observed as part of the legacy program \citep{meyer04}, DP Tau observed with the spectrograph at Gemini North \citep{herczeg07}, and T Tau N and S observed with the spectrograph at the Very Large Telescope (VLT) \citep{boekel09}."
. These references describe the respective data reduction in detail., These references describe the respective data reduction in detail.
 Our targets and their properties are listed in Tables 1-—4 (the first ten entries are displayed: the complete tables are available in the electronic version of thispaper)., Our targets and their properties are listed in Tables \ref{table1}- \ref{table4} (the first ten entries are displayed; the complete tables are available in the electronic version of thispaper).
 Table | lists the adopted stellar names. the same as those used in the," Table \ref{table1} lists the adopted stellar names, the same as those used in the"
Th anv respects. one of the most extraordinary globular clusters within our Milky Way isContam.,"In many respects, one of the most extraordinary globular clusters within our Milky Way is."
 It is uot ouly the most massive (3ον108 AL.) aud with a projected cllipticity of 0.18 one of the most flattened clusters in our galaxy. but it furthermore has a retrograde orbit around the Calactic centre. unlike most other globular clusters.," It is not only the most massive $3\times 10^{6}\; M_{\odot}$ ) and with a projected ellipticity of 0.18 one of the most flattened clusters in our galaxy, but it furthermore has a retrograde orbit around the Galactic centre, unlike most other globular clusters."
" Although lis one of the lougest-known- globular clusters and has been the subject of countless studies, the history of its oewestigation is more oue of adding chigmiatic properties than explaining them,"," Although is one of the longest-known globular clusters and has been the subject of countless studies, the history of its investigation is more one of adding enigmatic properties than explaining them."
 Already ? realised the tumsnal width of the red giant branch (RGB). which was later explained as a spread in metallicity by ¢ and which was spectroscopically coufirined by ?..," Already \citet{D/W:67} realised the unusual width of the red giant branch (RGB), which was later explained as a spread in metallicity by \citet{C/S:73} and which was spectroscopically confirmed by \cite{F/R:75}."
 Further spectroscopic analyses revealed significant variationsρα in nearly all (ciment abundances (e.g.22).," Further spectroscopic analyses revealed significant variations in nearly all element abundances \citep[e.g.][]{N/DC:95, S/S:00}."
 Especially remarkable is the wide spread in |Fe/H]. which covers a range of 3140 to 0.5dex.," Especially remarkable is the wide spread in [Fe/H], which covers a range of $-2.0$ to $-0.5\unit{dex}$."
 Both photometric and spectroscopic studies point to multiple stellar populations withinCon. which differ not only iu metal abundances but also in their spatial aud kinematic properties.," Both photometric and spectroscopic studies point to multiple stellar populations within, which differ not only in metal abundances but also in their spatial and kinematic properties."
 The stars on the RGB iu cean be divided into at least 3 supopulations., The stars on the RGB in can be divided into at least 3 sub-populations.
hat 4N dominant metal-poor compoucut (peaking at |[Fe/TI|~L.7 dex) contributes about of the stars., A dominant metal-poor component (peaking at $\sim -1.7\unit{dex}$ ) contributes about of the stars.
 This population rotates at a maximum velocity of τονωςSkinlas (?)..7).. Iu contrast.Ontrast. thee," This population rotates at a maximum velocity of $V_{\rm rot,max} \sim 8\unit{km \: s^{-1}}$ \citep{M/M/M:97}."
" intermediateintermeciate metallicity popukxopulatioulati E(about1 of thed starsi iand1 peasnmupeaking"" iatf Πιν—L2 dex) shows no rotation but rather a concentrationQ( towardsQWCLS the( chisterSÍC centre( (7)."," In contrast, the intermediate metallicity population (about of the stars and peaking at $\sim 
-1.2\unit{dex}$ ) shows no rotation but rather a concentration towards the cluster centre \citep{N/F:97}."
de The( inostO8 metal-rich compoucut courprises about of the stars with iuetallicities around 0.6dex. aud its spatial distribution is off-ceutre with respect to the more »)or populatious (?2)..," The most metal-rich component comprises about of the stars with metallicities around $-0.6\unit{dex}$, and its spatial distribution is off-centre with respect to the more metal-poor populations \citep{H/R:00}."
 One has to be aware of the changes ii nomenclature of these populations over the vears., One has to be aware of the changes in nomenclature of these populations over the years.
 Before the discovery of the most ietal-rich population (77). what are here called metal-poor aud intermediate qaetallicity populations are often referred to as metal-poor and metalrich.," Before the discovery of the most metal-rich population \citep{L/J:00, P/F:00}, what are here called metal-poor and intermediate metallicity populations are often referred to as metal-poor and metal-rich."
 Recent ligh-resolution umulti-baud images reveal that tye population puzzle in i; probably even more complicated., Recent high-resolution multi-band images reveal that the population puzzle in is probably even more complicated.
 2 discovered. apart your the main subgiaut brauch (SCB). an adcditional. uzurow. aud well-defined SCB in the colom-magnitude diagram (CMD) that secius to be the extension of the very sq RGB found by ? aud ?..," \citet{F/S:04}
 discovered, apart from the main subgiant branch (SGB), an additional, narrow, and well-defined SGB in the colour-magnitude diagram (CMD) that seems to be the extension of the very red RGB found by \citet{L/J:00} and \citet{P/F:00}."
 Also. finer substrietures have Ίσσα identified ou the RGB aud the main sequence (MS) qune DIST photometry.," Also, finer substructures have been identified on the RGB and the main sequence (MS) using HST photometry."
 ? have revealed the subdivision of (bo intermediate mctalliciv population on the RGB inte at east 3 discrete populations., \citet{S/F/P:05} have revealed the subdivision of the intermediate metallicity population on the RGB into at least 3 discrete populations.
 Very stunning is also the ification of the MS (ee.?).., Very stunning is also the bifurcation of the MS \citep[e.g.][]{B/P:04}.
 The observed fact that the yed dranch of the MS contains the majority of the stars. which is the opposite of what is observed on the RGB. seas at first very puzzling.," The observed fact that the red branch of the MS contains the majority of the stars, which is the opposite of what is observed on the RGB, was at first very puzzling."
 An unusual elim euxiclinieut of the intermediate aictallicity population now seen to pc one of the most favoured explanations for this feature (oco.22).," An unusual Helium enrichment of the intermediate metallicity population now seems to be one of the most favoured explanations for this feature \citep[e.g.][]{N:04, P/V:05}."
 Both spectroscopic and broad and narrow band photometric studies sugeest a spread im age (e.g.22777).," Both spectroscopic and broad and narrow band photometric studies suggest a spread in age \citep[e.g.][]{L/J:00, H/W:00, H/R:00, R/L:04, H/W:04}."
 Allottjose independent imeasureimeuts reveal a siguificaut aee difference between the sub-populatious of 2-6 Civi. correlated in the seuse that the vouuger populations are the more metalrich ones.," All of these independent measurements reveal a significant age difference between the sub-populations of 2-6 Gyr, correlated in the sense that the younger populations are the more metal-rich ones."
 Neither- the |ospread dn |Fe/II| nor iu age has been observed iu: anv other elobular. cluster so far., Neither the spread in [Fe/H] nor in age has been observed in any other globular cluster so far.
: Presinablx. μις gynation history of deiffersHag fundamentallyd +from the ones oft the other globular clusters in the Milkv Was.," Presumably, the formation history of differs fundamentally from the ones of the other globular clusters in the Milky Way."
 There are many different possibilities discussed as conceivable origins of the variatious in metal abundances aud the observed age spread., There are many different possibilities discussed as conceivable origins of the variations in metal abundances and the observed age spread.
 One of the most convincing ideas is that, One of the most convincing ideas is that
Fig.,Fig.
 | shows the solutions as a function of dimensionless period 1/5., 1 shows the solutions as a function of dimensionless period $1/\beta$.
 The upper branch shows that. rapidly BE black holes plus accretion disk form in small-period binaries (Bethe.Brown&Lee2003:Lee.Wijers2002).. following a surge for periods bevond the bifurcation points The resultingo mass and energyo lractions as a finetion of 1/2/ are shown in Fig.," The upper branch shows that rapidly spinning black holes plus accretion disk form in small-period binaries \citep{bet03,lee02}, following a surge for periods beyond the bifurcation points The resulting mass and energy fractions as a function of $1/\beta$ are shown in Fig."
o 2., 2.
 These {wo egeometries serve to bound the rangee of values in more detailed caleulations. e.g..e throughe multi-dimensional numerical simulations.," These two geometries serve to bound the range of values in more detailed calculations, e.g., through multi-dimensional numerical simulations."
trajectory corresponds to continume accretion bevond surge. wherein matter remaining in the remnant envelope forms an accretion disk outside the ISCO.," corresponds to continuing accretion beyond surge, wherein matter remaining in the remnant envelope forms an accretion disk outside the ISCO."
 At this point. magnetohvdrodyvnamical stresses within the disk as well as disk winds may drive continuing accretion.," At this point, magnetohydrodynamical stresses within the disk as well as disk winds may drive continuing accretion."
 Accretion from the ISCO onto the black hole further increases the black hole mass and spin according to 2M?=const. (Dardeen1970).. generally causing spin-up towards an extremal state of the black hole.," Accretion from the ISCO onto the black hole further increases the black hole mass and spin according to $zM^2=\mbox{const.}$ \citep{bar70}, generally causing spin-up towards an extremal state of the black hole."
 In Fig., In Fig.
 1 this is indicated by accretion beyond the upper ISCO-branch., 1 this is indicated by accretion beyond the upper ISCO-branch.
"spin-down corresponds to a lone-duration burst of gravitational radiation emitted by a non-anxisvinmetric torus (vanPutten1999;vanetal.2004).. described bv a frequency ancl energy where Mz=AL/TAL. and ji—AMp/0.03AM and ap=Oy/Qy, denote the relative mass and angular velocity of the torus."," corresponds to a long-duration burst of gravitational radiation emitted by a non-axisymmetric torus \citep{mvp99,mvp04}, described by a frequency and energy where $M_7=M/7M_\odot$, and $\mu=M_T/0.03M$ and $\eta=\Omega_T/\Omega_H$ denote the relative mass and angular velocity of the torus."
 This takes place if the torus is uniformly magnetized with the remnant magnetic field of (he progenitor star., This takes place if the torus is uniformly magnetized with the remnant magnetic field of the progenitor star.
 In Fig., In Fig.
 1. this radiative spin-dowun is indicated by a transition from (he upper ISCO branch to the branch on which the angular velocities of the black hole and of matter at the ISCO match (Ὁμ=Ojsco and 5= 1).," 1, this radiative spin-down is indicated by a transition from the upper ISCO branch to the branch on which the angular velocities of the black hole and of matter at the ISCO match $\Omega_H=\Omega_{ISCO}$ and $\eta=1$ )."
 This radiative transition lasts for the lifetime of rapid DA of the black hole — a dissipative timescale of tens of seconds (vanPutten&Levinsona2003).., This radiative transition lasts for the lifetime of rapid spin of the black hole – a dissipative timescale of tens of seconds \citep{mvp03a}.
 Additional matter accreted is either blown off the torus in its winds. or accumulates accretes onto the black hole after spin-down.," Additional matter accreted is either blown off the torus in its winds, or accumulates and accretes onto the black hole after spin-down."
"Precise age estimates of high redshift (5) galaxies directly constrain the epoch of galaxy formation.z,.. where is defined as the epoch when the majority of stars formed.","	Precise age estimates of high redshift $z$ ) galaxies directly constrain the epoch of galaxy formation, where is defined as the epoch when the majority of stars formed."
 Constraming τεῖς important to cosmology., 	Constraining is important to cosmology.
 For example. one of the key questions in modern cosmology has been whether the majority of stars in giant elliptical galaxies form at high redshifts through violent starbursts or during rather recent merger/interaction activities between smaller galaxies.," 	For example, one of the key questions in modern cosmology has been whether the majority of stars in giant elliptical galaxies form at high redshifts through violent starbursts or during rather recent merger/interaction activities between smaller galaxies."
" In addition. the age of a galaxy is a unique product of just a few cosmological parameters (e.g.. ο, A.Παν ος thus. it can be used to constrain cosmological parameters as well."," 	In addition, the age of a galaxy is a unique product of just a few cosmological parameters (e.g., $\Omega$, $\Lambda$, ); thus, it can be used to constrain cosmological parameters as well."
 Spinrad and his collaborators recently obtained the rest-frame UV spectrum of 553W091. a very red galaxy at z— 1.552. using the Keck Telescope (Dunlop et al.," 	Spinrad and his collaborators recently obtained the rest-frame UV spectrum of 53W091, a very red galaxy at $z = 1.552$ , using the Keck Telescope (Dunlop et al."
 1996; Spinrad et al., 1996; Spinrad et al.
 1997)., 1997).
 Based on their analysis on the UV spectrum and A—K color. they have concluded that 553W09] is at least 3.5 Gyr old already at z=1.552. which suggests Ho< 45 km sec Μρο. in the context of the Einstein-de Sitter cosmology.," 	Based on their analysis on the UV spectrum and $R-K$ color, they have concluded that 53W091 is at least 3.5 Gyr old already at $z = 1.552$, which suggests $<$ 45 km $^{-1}$ $^{-1}$, in the context of the Einstein-de Sitter cosmology."
 !When more recently measured cosmological parameters are used (e.g.. Ho= 65. 0=0.3. A 20.7; Aldering et al.," 	When more recently measured cosmological parameters are used (e.g., = 65, $\Omega = 0.3$, $\Lambda = 0.7$ ; Aldering et al."
 1998). this age estimate suggests that LBDS 53W091 formed at zy>6.5.," 1998), this age estimate suggests that LBDS 53W091 formed at $z_{f} \gtrsim 6.5$."
 Spinrad et al, 	Spinrad et al.
s results have been disputed by two independent studies.,'s results have been disputed by two independent studies.
 Bruzual Magris (19972) combined the UV spectrum studied by Spinrad et al., 	Bruzual Magris (1997a) combined the UV spectrum studied by Spinrad et al.
 with R. J. H. K photometry and obtained an age estimate 1.4 Gyr.," with $R$, $J$, $H$, $K$ photometry and obtained an age estimate 1.4 Gyr."
 Heap et al. (, 	Heap et al. (
1998) interpreted the UV spectral breaks using model atmosphere results specifically constructed for their study and using the 1997 version Yale tsochrones with no convective core overshoot (see $3).,1998) interpreted the UV spectral breaks using model atmosphere results specifically constructed for their study and using the 1997 version Yale isochrones with no convective core overshoot (see 3).
 They have estimated that the age of 553W09] lies between | and 2 Gyr., They have estimated that the age of 53W091 lies between 1 and 2 Gyr.
 They found that the Kurucz spectral library. which ts currently used in virtually all population synthesis studies including that of Spinrad et al.," 	They found that the Kurucz spectral library, which is currently used in virtually all population synthesis studies including that of Spinrad et al.,"
 does not match the detailed spectral features in the UV spectrum of an F-type main sequence (MS) star obtained with HST/STIS. in terms of the magnitudes of the spectral breaks used in the analysis of Spinrad et al. (," does not match the detailed spectral features in the UV spectrum of an F-type main sequence (MS) star obtained with HST/STIS, in terms of the magnitudes of the spectral breaks used in the analysis of Spinrad et al. ("
1997).,1997).
 This implies that Spinrad et al, 	This implies that Spinrad et al.
/s age estimate using UV spectral breaks may suffer from some systematic errors.,'s age estimate using UV spectral breaks may suffer from some systematic errors.
 Such discordant age estimates undermine our efforts to use this important technique as a probe of cosmology., Such discordant age estimates undermine our efforts to use this important technique as a probe of cosmology.
 We have carried out a similar exercise. estimating the age of this galaxy. using the Yi population synthesis models (Y1. Demarque. Oemler 1997).," 	We have carried out a similar exercise, estimating the age of this galaxy, using the Yi population synthesis models (Yi, Demarque, Oemler 1997)."
 In this paper. we present the results from the UV — visible continuum analysis only.," 	In this paper, we present the results from the UV – visible continuum analysis only."
 An analysis on the UV spectral breaks is currently in. progress., 	An analysis on the UV spectral breaks is currently in progress.
 We attemptto improve our age estimate by adopting (1), 	We attemptto improve our age estimate by adopting (1)
challenge in models of the cud stage of planet formation.,challenge in models of the end stage of planet formation.
 The dyvnaiical interactions between the protoplauets uust be calculated by direct N-body techuiques. but the sanaller fragments will need to be treated statistically.," The dynamical interactions between the protoplanets must be calculated by direct $N$ -body techniques, but the smaller fragments will need to be treated statistically."
 As a result. detailed models of the giant impact stage uust adopt new aethods. such as the hwbrid codes nentioncd above.," As a result, detailed models of the giant impact stage must adopt new methods, such as the hybrid codes mentioned above."
" Ivbrid mocels are needed to be able o calculate what fraction of the debris is reaccreted outo he final planets aud what fraction 1s removed (ο.ο,, via Povuting-Robertsou drag)."," Hybrid models are needed to be able to calculate what fraction of the debris is reaccreted onto the final planets and what fraction is removed (e.g., via Poynting-Robertson drag)."
 For fragments ground down low about 1l ian in size. new disruption criteria are still needed in the streneth regiae that fully account for uaterial properties. pact angle. mass ratio. aud mipact velocity.," For fragments ground down below about 1 km in size, new disruption criteria are still needed in the strength regime that fully account for material properties, impact angle, mass ratio, and impact velocity."
 The influence of more realistic collision plysics on the time scale of terrestrial planet formation is difficult to predict because of compcting factors., The influence of more realistic collision physics on the time scale of terrestrial planet formation is difficult to predict because of competing factors.
 Ou one haul. the growth rate of planets is slower when outcomes other than perfect mereine are inclided.," On one hand, the growth rate of planets is slower when outcomes other than perfect merging are included."
 Ou the other hand. the debris produced by collisions can influence the overall dvnamics ofthe embryos and planctesimals (6.8.. via dvunamical friction. which could lead to more lower velocity collisions and more frequent merging outcomes).," On the other hand, the debris produced by collisions can influence the overall dynamics of the embryos and planetesimals (e.g., via dynamical friction, which could lead to more lower velocity collisions and more frequent merging outcomes)."
 The time scale for the cud stage of planct formation was investigated in a recent study that included two collision outcomes., The time scale for the end stage of planet formation was investigated in a recent study that included two collision outcomes.
 conducted N-body simulations with a collision model that allowed or either perfect iiereiug or an ideal hit-and-runu eveut using their enpirical boundary presented in equation 15 Qvlich they applied at all impact angles)., conducted $N$ -body simulations with a collision model that allowed for either perfect merging or an ideal hit-and-run event using their empirical boundary presented in equation \ref{eqn:vhr} (which they applied at all impact angles).
 In a hit-audam collision. neither body lost mass (there was no yagiuentation) but the relative velocities of the bodies decreased.," In a hit-and-run collision, neither body lost mass (there was no fragmentation) but the relative velocities of the bodies decreased."
 Iu simulations that began with 16 equaluass clubrves between 0.5 aud 1.5 AU. they found that about half of the collisions were hit-aud-run events.," In simulations that began with 16 equal-mass embryos between 0.5 and 1.5 AU, they found that about half of the collisions were hit-and-run events."
 After a hit-ancdaun eucounuter. the reduced relative velocities ed to a lieh probability of mergiug on the subsequent πο," After a hit-and-run encounter, the reduced relative velocities led to a high probability of merging on the subsequent encounter."
 As a result. the time scale for planet growth was esseutiallv the same as in simulations with the same starting conditions that assumed perfect mereiue.," As a result, the time scale for planet growth was essentially the same as in simulations with the same starting conditions that assumed perfect merging."
 The magnitude of any changes to plauct erowtl time scales frou different collisiou outcomes is sensitive to the initial conditions in the simulation., The magnitude of any changes to planet growth time scales from different collision outcomes is sensitive to the initial conditions in the simulation.
" The simulations by began with few relatively large embryos (0.1537,.) aud no planetesinals.", The simulations by began with few relatively large embryos $0.15M_{\Earth}$ ) and no planetesimals.
" With +us initial distribution of e1ibrvos. all collisions were between comparable mass bodies CAZ,/M,> 0.1). which are the most likely to be graze-audanerse or hit-aud-run events (Figure 5))."," With this initial distribution of embryos, all collisions were between comparable mass bodies $M_{\rm p}/M_{\rm t}>0.1$ ), which are the most likely to be graze-and-merge or hit-and-run events (Figure \ref{fig:collmapequal}) )."
 Since the hit-aud-aunu bodies are likely to recollide ou time scales comparable to the orbital period. the overall effect of hit-and-run evenuts ou this very cud stage of planet erowth was found to be negligible.," Since the hit-and-run bodies are likely to recollide on time scales comparable to the orbital period, the overall effect of hit-and-run events on this very end stage of planet growth was found to be negligible."
 If fragmentation were included or if the sizes of the bodies were more diverse (e.g. with the inclusion of plauctesimals or siualler initial enübirvos). more realistic collision outcomes would have a larecr effect on planct growth.," If fragmentation were included or if the sizes of the bodies were more diverse (e.g., with the inclusion of planetesimals or smaller initial embryos), more realistic collision outcomes would have a larger effect on planet growth."
 Asx mentioned above. using the sale initial conditions in their multi-scale calculations. found that fracmenutation was a sjenificaut process.," As mentioned above, using the same initial conditions in their multi-scale calculations, found that fragmentation was a significant process."
 Although fragmentation makes planet erowth from au individual collision less cficient. other effects from the debris mav lead to faster plauct growth overall or faster stages of planet erowth.," Although fragmentation makes planet growth from an individual collision less efficient, other effects from the debris may lead to faster planet growth overall or faster stages of planet growth."
 For example. during oligarchic erowth. found that fragmentation led o faster erowth of eimibrvos because smaller fragments are maiore easilv captured.," For example, during oligarchic growth, found that fragmentation led to faster growth of embryos because smaller fragments are more easily captured."
 However. fragnieutatiou also decreased the surface density of solids in the disk (which nuits the enmibrvo'« final mass) because fragments were ost nore quickly bv drag processes as they were erol down in size.," However, fragmentation also decreased the surface density of solids in the disk (which limits the embryo's final mass) because fragments were lost more quickly by drag processes as they were ground down in size."
" Daving the eiut nuact stage, forud that the strongo dynamical fiction roni 1000. plauctesimals led to overall faster planet erowth compares to studies withou any planetesinials."," During the giant impact stage, found that the strong dynamical friction from 1000 planetesimals led to overall faster planet growth compared to studies without any planetesimals."
 Ilowever. them sudy assumed perect merging for a1 collision outcomes and non-interacting planctesimals.," However, their study assumed perfect merging for all collision outcomes and non-interacting planetesimals."
 If dvnaiical friction becomes very large. a gap could form in the planetesimal disk around au cmbirvoe. which would effectively halt the erowth of that emibrvo.," If dynamical friction becomes very large, a gap could form in the planetesimal disk around an embryo, which would effectively halt the growth of that embryo."
 ence. the role of fragmentation on planet erowth time scales is not independent of other processes acting at the same time.," Hence, the role of fragmentation on planet growth time scales is not independent of other processes acting at the same time."
" Ultimately, because fragmentation is a critical process that feeds back into the dynamics of plauct growth. new simulations that include both eumbryvos aud a fully interacting population of small bodies are needed to investigate how amore realistic collision outcomes affect formation time scales."," Ultimately, because fragmentation is a critical process that feeds back into the dynamics of planet growth, new simulations that include both embryos and a fully interacting population of small bodies are needed to investigate how more realistic collision outcomes affect formation time scales."
 We found that the core-to-meautle lnass fraction increases duiug the erowth of planets via fragmicutation during collimous between differcutiaed oemibrvos., We found that the core-to-mantle mass fraction increases during the growth of planets via fragmentation during collisions between differentiated embryos.
 Loss of the silicate mautle primarily occurred during partial accretion aud erosion of the projecie in hit-aud-runu events., Loss of the silicate mantle primarily occurred during partial accretion and erosion of the projectile in hit-and-run events.
 Frou our group D. Moute Caro caleulatious. the mean merease in the core mass fracticπι was for the 3l lareest plaucts (Figure TEE. blacς histogram).," From our group B Monte Carlo calculations, the mean increase in the core mass fraction was for the 34 largest planets (Figure \ref{fig:montecarlo}E E, black histogram)."
 The nagnitude of the core fraction increase is poteutially observable iu the study of the chemical composition of ueuets aud early Solar System materials., The magnitude of the core fraction increase is potentially observable in the study of the chemical composition of planets and early Solar System materials.
 argue that the Earth's mlk irou to magnesium (Fe/Me)— ratio is siguificautlv arecr than the solu ratio., argue that the Earth's bulk iron to magnesium (Fe/Mg) ratio is significantly larger than the solar ratio.
 They propose that the Earth lost a portion of its silicate mantle durius the eiut impact stage. which raised the Fe/Ale ratio of he final planet compared to the more primitive (closer o nebular composition)- imaterials that formed the Xauetarv enmibrvos.," They propose that the Earth lost a portion of its silicate mantle during the giant impact stage, which raised the Fe/Mg ratio of the final planet compared to the more primitive (closer to nebular composition) materials that formed the planetary embryos."
 estimate hat the whole-Earth Fe/Ale amass ratio is 2.1+ L1., estimate that the whole-Earth Fe/Mg mass ratio is $2.1\pm0.1$ .
 The Fe/Ale value for primitive materials is not shown precisely., The Fe/Mg value for primitive materials is not known precisely.
 The solar plotosphere value. 1.570.2005).. is too poorly constrained toL vc useful for such detailed comparisons.," The solar photosphere value, $1.87\pm0.4$, is too poorly constrained to be useful for such detailed comparisons."
 The Fe/Mg ratio for the solar wind will be constrained bv Conuesis uissjon., The Fe/Mg ratio for the solar wind will be constrained by Genesis mission.
 Early results sueeested a lower value than he solar photosphere2011): however. final data calibration is still iu progress JJurewiez. cconuu.).," Early results suggested a lower value than the solar photosphere; however, final data calibration is still in progress Jurewicz, comm.)."
 The Fe/Mes ratio for carbouaceous choudrites. the most primitive type of meteorite. is 1.92+40.0820051.," The Fe/Mg ratio for carbonaceous chondrites, the most primitive type of meteorite, is $1.92\pm0.08$."
 The available data sugeest that the Earth is eunrched iu Fe/Ale compared to solar composition bv approximatcly., The available data suggest that the Earth is enriched in Fe/Mg compared to solar composition by approximately.
. Our Monte Carlo calculations of the cumulative effects of realistic collision outcomes are in excelleut agreciuieut with the collisional erosion idea proposed by, Our Monte Carlo calculations of the cumulative effects of realistic collision outcomes are in excellent agreement with the collisional erosion idea proposed by
"However we note that the PAH feature in the GRB host is substantially weaker than in 77714 and 55461, while the main other PAHs detected atum,, and have roughly the same relative strengths in the three spectra.","However we note that the PAH feature in the GRB host is substantially weaker than in 7714 and 5461, while the main other PAHs detected at, and have roughly the same relative strengths in the three spectra."
" Furthermore, the mid-IR spectral slopes are almost identical between the WR region and 77714, while the mid-IR spectrum of the HII region 55461 seems to rise even more rapidly at the longest wavelengths."," Furthermore, the mid-IR spectral slopes are almost identical between the WR region and 7714, while the mid-IR spectrum of the HII region 5461 seems to rise even more rapidly at the longest wavelengths."
" Finally the IRS spectrum of the WR region reveals the presence of prominent fine-structure emission lines such asum,, andjum,, which[SIV]- are also detected in 55461 but either absent or much fainter in 77714."," Finally the IRS spectrum of the WR region reveals the presence of prominent fine-structure emission lines such as, and, which are also detected in 5461 but either absent or much fainter in 7714."
 In the following sections we will quantitatively analyze the relative strengths of the different features characterizing the WR region of the 9980425 host., In the following sections we will quantitatively analyze the relative strengths of the different features characterizing the WR region of the 980425 host.
 In particular we will further explore how it differs from what has been so far observed in the mid-IR. spectra of star-forming galaxies., In particular we will further explore how it differs from what has been so far observed in the mid-IR spectra of star-forming galaxies.
 .2cm This section describes the broad-band photometry measurements obtained from our MIPS imaging., .2cm This section describes the broad-band photometry measurements obtained from our MIPS imaging.
 All the measured fluxes are summarized in, All the measured fluxes are summarized in
Simple power law fits to the flux binned spectra from the monitoring observations after removal of the torus component still reveal a strong correlation of the spectral slope with the flux level of NGC 4051: the softest spectra are found when the flux is highest (see tig 7).,Simple power law fits to the flux binned spectra from the monitoring observations after removal of the torus component still reveal a strong correlation of the spectral slope with the flux level of NGC 4051: the softest spectra are found when the flux is highest (see fig \ref{lineflux}) ).
 In the context of the standard models for the X- emission of Seyfert galaxies three principal mechanisms can contribute to this correlation: In order to distinguish between the first two possibilities we titted models with a variable Gaussian line in addition © the fixed torus reflection model to the flux selected spectra of both the monitoring and December 1996 observations., In the context of the standard models for the X-ray emission of Seyfert galaxies three principal mechanisms can contribute to this correlation: In order to distinguish between the first two possibilities we fitted models with a variable Gaussian line in addition to the fixed torus reflection model to the flux selected spectra of both the monitoring and December 1996 observations.
 We reduced he number of free parameters by fixing some of the less critical owameters of the model to reasonable values., We reduced the number of free parameters by fixing some of the less critical parameters of the model to reasonable values.
 The ugh energy cutoff of the primary power law was set to 300 keV. effectively eliminating the cutoff.," The high energy cutoff of the primary power law was set to 300 keV, effectively eliminating the cutoff."
 The iron and metal abundances were set to the solar values and the inclination angle was fixed at 30 as suggested by the results of disk line model tits to ASCA spectra (Nandraetal.1997)., The iron and metal abundances were set to the solar values and the inclination angle was fixed at $30^\circ$ as suggested by the results of disk line model fits to ASCA spectra \cite{Nandra97}.
. Confidence contours of the primary photon index Land the reflected fraction /? are shown in Figs 3 and +., Confidence contours of the primary photon index $\Gamma$ and the reflected fraction $R$ are shown in Figs \ref{mon_cont} and \ref{dec_cont}.
 It is evident that a continuum of variable spectral index is needed to satisfy the data., It is evident that a continuum of variable spectral index is needed to satisfy the data.
 The best fit photon indices range from L=1.6 for the lowest flux states to --2.3 for the highest flux states of the monitoring observations., The best fit photon indices range from $\Gamma=1.6$ for the lowest flux states to $\Gamma=2.3$ for the highest flux states of the monitoring observations.
 Note that the inclusion of he presumed constant torus component does not change the results qualitatively., Note that the inclusion of the presumed constant torus component does not change the results qualitatively.
 Tf we omit this fixed component. the photon index of the lowest flux spectrum is I=1.2: the fit to the highest flux spectrum is not affected by this change.," If we omit this fixed component, the photon index of the lowest flux spectrum is $\Gamma=1.2$; the fit to the highest flux spectrum is not affected by this change."
 No variability of he reflected fraction is obvious and at the 004. confidence levels all spectra are consisistent with zero reflection., No variability of the reflected fraction is obvious and at the $90\%$ confidence levels all spectra are consisistent with zero reflection.
 Apart from the ‘act that the constant reflected (torus) component will become less prominent at higher flux levels. there is no obvious contribution of reflection to the spectral variability.," Apart from the fact that the constant reflected (torus) component will become less prominent at higher flux levels, there is no obvious contribution of reflection to the spectral variability."
 We also investigated the warm absorber as the third possible source of hardness variability., We also investigated the warm absorber as the third possible source of hardness variability.
 Again including the low state spectrum as a constant component. we modelled the flux-selected X-ray spectra with a power law and Gaussian emission line absorbed by a warm absorber model) with a column density of Vy=5-1077cm27. as suggested by Komossa&Fink (1999).," Again including the low state spectrum as a constant component, we modelled the flux-selected X-ray spectra with a power law and Gaussian emission line absorbed by a warm absorber model) with a column density of $N_{\rm H} =
5\cdot10^{22} {\rm cm^{-2}}$, as suggested by \scite{Komossa}."
. The models were calculated on a grid in spectral index Land ionisation parameter € and fitted to each of the 7 flux selected spectra of the monitoring observations., The models were calculated on a grid in spectral index $\Gamma$ and ionisation parameter $\xi$ and fitted to each of the 7 flux selected spectra of the monitoring observations.
 As a result. the ionisation parameter appears to be poorly constrained.," As a result, the ionisation parameter appears to be poorly constrained."
 The spectral index is clearly correlated with flux and there is no DL that would be consistent with all of the 7 spectra (see Table 4»)., The spectral index is clearly correlated with flux and there is no $\Gamma$ that would be consistent with all of the 7 spectra (see Table \ref{warmabs}) ).
 We therefore conclude that changes in the ionisation state of the warm absorber in NGC 4051 do not play a significant role in the spectral variability of the source., We therefore conclude that changes in the ionisation state of the warm absorber in NGC 4051 do not play a significant role in the spectral variability of the source.
 To summarize. variability of primary power law continuum slope is the only viable explanation of the strong hardness variability in the 2-24 keV X-ray spectrum of NGC 4051.," To summarize, variability of primary power law continuum slope is the only viable explanation of the strong hardness variability in the 2-24 keV X-ray spectrum of NGC 4051."
 After subtraction of the low state reflection spectrum with the narrow iron fluorescence line. broad fluorescence line emission is still detectable at all flux levels.," After subtraction of the low state reflection spectrum with the narrow iron fluorescence line, broad fluorescence line emission is still detectable at all flux levels."
 Figs., Figs.
 5. and 6. show ratio plots obtained by fitting power law models to the flux selected spectra," \ref{monline} and \ref{decline}
 show ratio plots obtained by fitting power law models to the flux selected spectra"
We included in our STIS target list three objects that were previously observed. with the Plauctary Camera Land the Faint Object Camera.,We included in our STIS target list three objects that were previously observed with the Planetary Camera 1 and the Faint Object Camera.
 The rationale teyduchide a few repeat targets from the sample of Stanehellinietal.(1999) was to check the morphological types and dimensions with both methods aud to asses the reliability of the pre-COSTAR archival data., The rationale to include a few repeat targets from the sample of \citet{sta99} was to check the morphological types and dimensions with both methods and to asses the reliability of the pre-COSTAR archival data.
 All three were observed before the firstLEST servicing uussion. thus the images suffer for the uncorrected προcal aberration of the telescope mirror.," All three were observed before the first servicing mission, thus the images suffer for the uncorrected spherical aberration of the telescope mirror."
 Iu auZ/ST archival study (Stanghellinietal.1999) it was found that iiorphology was well determined with the archived instrtunents. while dimensions of the nebulae with the photometric method were not always reliable because of the possible presence of fainter extended halos around the images.," In an archival study \citep{sta99} it was found that morphology was well determined with the archived instruments, while dimensions of the nebulae with the photometric method were not always reliable because of the possible presence of fainter extended halos around the images."
 Since proerani No., Since program No.
 8366 was a snapshot program. we could not euarautee exactly which targets would be observed for this comparison.," 8366 was a snapshot program, we could not guarantee exactly which targets would be observed for this comparison."
 Iu the eud. targets SAIP J UN 1). SMP 6 (N 6). aud SMP 22 (N 67) were with STIS.," In the end, targets SMP 1 (N 1), SMP 6 (N 6), and SMP 22 (N 67) were re-observed with STIS."
 The first two targets were observed with PC1 earlier. while SMP 22 had been observed with FOC.," The first two targets were observed with PC1 earlier, while SMP 22 had been observed with FOC."
 SMP 1 aud SAIP 6 are sinall nebulae with very uncomplicated structures;, SMP 1 and SMP 6 are small nebulae with very uncomplicated structures.
 The round (E?), The round (E?)
 morphology of these two targets was casily ποσα in the PCL data set (Vassiliadisetal.1998)., morphology of these two targets was easily seen in the PC1 data set \citep{vasea98}.
". The carly photometric radius measured for these two PNs were and0.152""... respectively for SMP 1 and SAIP 6 was close to our measurements of and0."," The early photometric radius measured for these two PNs were and, respectively for SMP 1 and SMP 6 was close to our measurements of and."
"19"". The situation of the bipolar planetary nebula SMP 22 is very differeut.", The situation of the bipolar planetary nebula SMP 22 is very different.
 The complete nebular iiorphliology hat we see in most emission lines of the STIS spectra. but in particular iu the lines. is not as evident in the FOC image (Stanehellinictal.1999).. Figure 5.," The complete nebular morphology that we see in most emission lines of the STIS spectra, but in particular in the lines, is not as evident in the FOC image \citep{sta99}, Figure 5."
 Furthermore. we mcasure a photometric racius of17.. while the radius from the FOC image is almost three times larger.," Furthermore, we measure a photometric radius of, while the radius from the FOC image is almost three times larger."
 We reaualyzed these measurements. aud noted that a radius of already included about 75:&( of the total flux. hit to encircle the shtA of he flux required radius.," We reanalyzed these measurements, and noted that a radius of already included about $\%$ of the total flux, but to encircle the $\%$ of the flux required radius."
 From the comparison of the old aud new SAIP 22 images we conclude that he main morphological featires of MC PNs are reliable from the pre-COSTAR imaees. even if the detailed norphology was not resolved.," From the comparison of the old and new SMP 22 images we conclude that the main morphological features of MC PNs are reliable from the pre-COSTAR images, even if the detailed morphology was not resolved."
 Ou the other haud. the measurements of the photomeric radi are not reliable.," On the other hand, the measurements of the photometric radii are not reliable."
 From our analysis we cau see that SAIC PNs are nicely. classified using the same morphological scheme as the LAIC and Calactic PNs., From our analysis we can see that SMC PNs are nicely classified using the same morphological scheme as the LMC and Galactic PNs.
 In order to increase the sample size for statistical purposes. we include in the present sample three additional SAIC) PNs described and classified. i Staughelliniet.al.(1999).," In order to increase the sample size for statistical purposes, we include in the present sample three additional SMC PNs described and classified in \citet{sta99}."
. After climinating unresolved objects aud repeats. we lave at our disposal a sample of 30 SAIC PNs whose 1101]xiologv is well deteziiued.," After eliminating unresolved objects and repeats, we have at our disposal a sample of 30 SMC PNs whose morphology is well determined."
" This SAIC PN sample constitute nearly 50 percent of a] known SAIC ΤΝ», lus we consider fairly representativo at least of the bright PNs."," This SMC PN sample constitute nearly 50 percent of all known SMC PNs, thus we consider fairly representative at least of the bright PNs."
 In Table 1 we give he statistics of PN 1101]phology for the SAIC sample as compared to the LMC and Calactic samples fro ushawetal.(2001)., In Table 4 we give the statistics of PN morphology for the SMC sample as compared to the LMC and Galactic samples from \citet{sha01}.
 The LAIC and SAIC samples in cols 2 aud 3 of Table 1. have )eecn selected i jicdnülar wavs. aud hey iive sinülar observational biases (but noue of the extreme selection biases tlat affect Calactic PN sunges toward the Galactic plane).," The LMC and SMC samples in columns 2 and 3 of Table 4, have been selected in similar ways, and they have similar observational biases (but none of the extreme selection biases that affect Galactic PN samples toward the Galactic plane)."
 While the fraction of round. PNs remains more or less the same in he two samples. the fractious of E and DC in the SAIC are respective votwice aix Lore-half those of the LMC.," While the fraction of round PNs remains more or less the same in the two samples, the fractions of E and BC in the SMC are respectively twice and one-half those of the LMC."
 The overall frequeney of αποο PN in the SMC is only sixty- percent ha of the LMC., The overall frequency of asymmetric PN in the SMC is only sixty- percent that of the LMC.
 This remarkable results stronelv sugecstsOO that the difference between the Magellanic Clouds is reflected in their PN populations: something iu the euvironmieut of the SAIC may not no be favorae to the formation of, This remarkable results strongly suggests that the difference between the Magellanic Clouds is reflected in their PN populations: something in the environment of the SMC may not not be favorable to the formation of
ALOS-2 (Turner 22001) and PN (Strüdder 22001) cameras were operating inMODE. using the filter.,"MOS-2 (Turner 2001) and PN (Strüdder 2001) cameras were operating in, using the filter."
 The optical/UV OM camera (Mason ct al., The optical/UV OM camera (Mason et al.
 2001) was operating inMODE. observing through the UVAI2 filter (20502450 A3).," 2001) was operating in, observing through the UVM2 filter (2050–2450 )."
 The RGS cameras were also in operation: however the count-rate was too low to allow phase-resolvecl analysis. which is the goal of this paper. thus they are only mentioned in passing here.," The RGS cameras were also in operation; however the count-rate was too low to allow phase-resolved analysis, which is the goal of this paper, thus they are only mentioned in passing here."
 Some aspects of these data were previously reported by Cropper ((2002)., Some aspects of these data were previously reported by Cropper (2002).
 We analysed the data using the software v5.4.1., We analysed the data using the software v5.4.1.
 The source data were extracted from a circular region of radius 3.5 aresee enclosing the source. with the backgroundd being5 taken from an annulus around this area.," The source data were extracted from a circular region of radius 3.5 arcsec enclosing the source, with the background being taken from an annulus around this area."
 Only single5 or double pixel events with a zero quality1 Dag5 were selected., Only single or double pixel events with a zero quality flag were selected.
 The N-rav and UV lighteurves are. presented. in Fig. L..," The X-ray and UV lightcurves are presented in Fig. \ref{fig:curve},"
 binned at 64 s ancl 32 s respectively., binned at 64 s and 32 s respectively.
 We show the sum of the MOS lighteurves: the PN lighteurve is very similar to the ALIOS curves and so is not shown., We show the sum of the MOS lightcurves; the PN lightcurve is very similar to the MOS curves and so is not shown.
 In any case. the Iickering level is higher than the photon noise.," In any case, the flickering level is higher than the photon noise."
 Power spectra of these lishteurves are given in Fig. 2.., Power spectra of these lightcurves are given in Fig. \ref{fig:ft}.
 The dominant pulses are the 4.55-hr and. 20.9-min orbital and spin pulses (we label the frequencies and rrespectively). and there is also power at the aand z|© sidebands.," The dominant pulses are the 4.85-hr and 20.9-min orbital and spin pulses (we label the frequencies and respectively), and there is also power at the and $\w+\W$ sidebands."
 Power at the beat period )) is indicative of spin interaction in the accretion process., Power at the beat period ) is indicative of spin--orbit interaction in the accretion process.
 N-ray beat periods are generally interpreted as resulting from accreting material which couples directly to the white cwarl magnetosphere. without passing through an accretion disc (e.g. streameofed accretion. or disc-overllow: accretion: Hellier 119892: Lellier 1991: Wynn IxIxing 1992).," X-ray beat periods are generally interpreted as resulting from accreting material which couples directly to the white dwarf magnetosphere, without passing through an accretion disc (e.g. stream-fed accretion, or disc-overflow accretion; Hellier 1989a; Hellier 1991; Wynn King 1992)."
 However. we see approximately equal power at the aand w|© sideband. frequencies. which may result. simply rom an amplitude modulation. of the spin pulse at. the orbital timescale (see Warner 1986). rather than from a stream.magnetosphere interaction.," However, we see approximately equal power at the and $\w+\W$ sideband frequencies, which may result simply from an amplitude modulation of the spin pulse at the orbital timescale (see Warner 1986), rather than from a stream–magnetosphere interaction."
 To test this we used the ieht-curve reconstruction program of 193 to see whether we could reproduce the light curve using a spin pulse multiplied » an orbital modulation. without any intrinsic modulation a thea »periocl.," To test this we used the light-curve reconstruction program of H93 to see whether we could reproduce the light curve using a spin pulse multiplied by an orbital modulation, without any intrinsic modulation at the period."
 The resulting mocel lishteurve is given in the bottom yanel of Fig. 1..," The resulting model lightcurve is given in the bottom panel of Fig. \ref{fig:curve},"
 and the power spectrum of the residuals to his fitted model is shown in the lower panel Fig. 2.., and the power spectrum of the residuals to this fitted model is shown in the lower panel Fig. \ref{fig:ft}.
 The absence of peaks significantly. above the flickering implies hat the model is adequate in reproducing the periodic components of the lishteurve., The absence of peaks significantly above the flickering implies that the model is adequate in reproducing the periodic components of the lightcurve.
 Thus we conclude that there was little or no disc-overllow accretion occurring at the time of the observations., Thus we conclude that there was little or no disc-overflow accretion occurring at the time of the observations.
 Note that. Hellier (1991) and. BOS reported. that the relative amplitude of the beat and spin pulses change with time. and suggested that the accretion mode of FO qr changes from periods with significant clisc-overllow accretion," Note that Hellier (1991) and B98 reported that the relative amplitude of the beat and spin pulses change with time, and suggested that the accretion mode of FO Aqr changes from periods with significant disc-overflow accretion"
 , 
because of the high-quality multiwavelength. information available for the 2deg? area of the Cosmic Evolution Survey(COSALOS: Scoville et al.,because of the high-quality multi-wavelength information available for the $2~\mathrm{deg}^2$ area of the Cosmic Evolution Survey; Scoville et al.
 2007) where the two systems were initially cetected in N-ravs. (Finoguenov. et al., 2007) where the two systems were initially detected in X-rays (Finoguenov et al.
 in preparation: G09)., in preparation; G09).
 In. particular. the identification as fossil and the selection. of the member-cancidate galaxies of either. group rest on robust photometric redshifts. and fapox22.5mag limited spectroscopy.," In particular, the identification as fossil and the selection of the member-candidate galaxies of either group rest on robust photometric redshifts, and $i_{\mathrm AB} \le 22.5~\mathrm{mag}$ limited spectroscopy."
 Interestingly. 1ο large-scale structure (LSS) in which the two fossil groups. are. embedded: is very different: | 0212.6 appears isolated. whereas | 0140.5. belongs to a LSS that covers the entire 2deg area of the (corresponding to à cross size of about 25.5Alpe at z5ε 0.4) and is traced by a total of 28 X-ray emitting groups.," Interestingly, the large-scale structure (LSS) in which the two fossil groups are embedded is very different: $+$ 0212.6 appears isolated, whereas $+$ 0140.8 belongs to a LSS that covers the entire $2~\mathrm{deg}^2$ area of the (corresponding to a cross size of about $25.5~\mathrm{Mpc}$ at $z \approx 0.4$ ) and is traced by a total of 28 X-ray emitting groups."
 Comparison ofthe GSME of these fossil groups and the composite GSME of the X-ray emitting groups at 0.3x20.5 that span a similar range in total mass (Ciodini et al., Comparison of the GSMF of these fossil groups and the composite GSMF of the X-ray emitting groups at $0.3 \le z \le 0.5$ that span a similar range in total mass (Giodini et al.
 in preparation) highlights a lack of star-forming galaxies with ALY107AL. in the (svo. fossil groups., in preparation) highlights a lack of star-forming galaxies with $M^{\mathrm{stars}} \ge 10^{10}~\mathrm{M}_{\sun}$ in the two fossil groups.
 Consistently. the distribution of their photometric member galaxies in the b rvs? colour.magnitude ciagranr reveals a well-defined red sequence and a Lack of significantly bluer. luminous galaxies oul to the virial radius. foo of either fossil group.," Consistently, the distribution of their photometric member galaxies in the $b - r$ vs $i$ colour–magnitude diagram reveals a well-defined red sequence and a lack of significantly bluer, luminous galaxies out to the virial radius $R_{200}$ of either fossil group."
 In addition. the total stellar mass fraction of these groups. within O.7Rooy is relatively large for their total masses (609). in particular for the less massive one Le. | 0212.6)," In addition, the total stellar mass fraction of these groups within $0.7 R_{200}$ is relatively large for their total masses (G09), in particular for the less massive one (i.e., $+$ 0212.6)."
 This evidence suggests that the overall conversion of (cold) gas into stars was accelerated and/or more ellicient in fossil groups with respect to the average X-ray emitting eroup in the same mass range and at similar redshift., This evidence suggests that the overall conversion of (cold) gas into stars was accelerated and/or more efficient in fossil groups with respect to the average X-ray emitting group in the same mass range and at similar redshift.
 The star-formation activity in member galaxies rapidly ceased afterwards., The star-formation activity in member galaxies rapidly ceased afterwards.
 At the same time. no significant infall of galaxies with APT>LOMALS took place in the last 3:6 Gyr for either fossil group at zz0.4.," At the same time, no significant infall of galaxies with $M^{stars} \ge 10^{10}~\mathrm{M}_{\sun}$ took place in the last 3–6 Gyr for either fossil group at $z \approx 0.4$."
 If the processes shaping the galaxy component of a fossil group seem to be independent of the LSS. the progenitor of an observed fossil group may depend: on it.," If the processes shaping the galaxy component of a fossil group seem to be independent of the LSS, the progenitor of an observed fossil group may depend on it."
 On the basis of the measured stellar. mass. fractions (6109) ancl X-rav-to-optical luminosity ratios (Ciodini et al., On the basis of the measured stellar mass fractions (G09) and X-ray-to-optical luminosity ratios (Giodini et al.
 in preparation). we propose that compact groups are the most likely progenitors of observed. low mass ancl isolated fossil groups. like | O212.¢x," in preparation), we propose that compact groups are the most likely progenitors of observed low mass and isolated fossil groups, like $+$ 0212.6."
 Conversely. the relatively carly infall of massive satellites likely originates massive fossil groups in dense regions of the LSS. like | 0140.5," Conversely, the relatively early infall of massive satellites likely originates massive fossil groups in dense regions of the LSS, like $+$ 0140.8."
 We thank the referee. Paul Eigenthaler. for a careful reacing of the original version of the paper ancl the useful comments that improved the presentation of the paper.," We thank the referee, Paul Eigenthaler, for a careful reading of the original version of the paper and the useful comments that improved the presentation of the paper."
 DP acknowledges the kind and fruitful hospitality at the [ürr extraterrestrisehe Physik (ALPE)., DP acknowledges the kind and fruitful hospitality at the Max-Planck-Institut fürr extraterrestrische Physik (MPE).
the example above. we would. like a mark that encodes information about which halos host which galaxies.,"the example above, we would like a mark that encodes information about which halos host which galaxies."
" One set of such marks would be functions of the local density. which can be computed in a number of wavs. e.g. the distance to the n""| nearest neighbor. the number of neighbors within a fixed metric aperture. spline kernel interpolation (e.g.Dehnen2001) or kernel deprojection (e.g.Eisenstein2003)."," One set of such marks would be functions of the local density, which can be computed in a number of ways, e.g. the distance to the $n^{\rm th}$ nearest neighbor, the number of neighbors within a fixed metric aperture, spline kernel interpolation \citep[e.g.][]{Deh01} or kernel deprojection \citep[e.g.][]{Eis03}."
. Massive halos tend to host more galaxies with higher density than lower mass halos. so if the HOD is changed we expect the densitv-marked correlation function to dilfer.," Massive halos tend to host more galaxies with higher density than lower mass halos, so if the HOD is changed we expect the density-marked correlation function to differ."
 The relation of local density to the number density of groups «r clusters (which are known to break degeneracies in. moclel fitting. e.g. Zheng&Weinberg 2007)) can be complex. but is casily calculable from a mock catalog.," The relation of local density to the number density of groups or clusters (which are known to break degeneracies in model fitting, e.g. \citealt{ZheWei07}) ) can be complex, but is easily calculable from a mock catalog."
 What function of p should we choose as our mark?, What function of $\rho$ should we choose as our mark?
" The choice is arbitrary. but. p""/(p7.|p"") has the nice property that it tends to zero for p«p, and unity for pX p,. the rapidity of the transition being controlled by à»."," The choice is arbitrary, but $\rho^n/(\rho_\star^n+\rho^n)$ has the nice property that it tends to zero for $\rho\ll\rho_\star$ and unity for $\rho\gg\rho_\star$ , the rapidity of the transition being controlled by $n$."
 This means the dynamic range in the mark is limited. which leads to more stable results.," This means the dynamic range in the mark is limited, which leads to more stable results."
 Lowe are concerned that our density estimator may be noisy. which is often true in practice. we should choose a low value of 5.," If we are concerned that our density estimator may be noisy, which is often true in practice, we should choose a low value of $n$."
 Hereafter we choose n=I., Hereafter we choose $n=1$.
 We now measure the local-clensity marked correlation 'unction for our two examples LLODs., We now measure the local-density marked correlation function for our two examples HODs.
 To begin we imagine hat we can use spectroscopy or multi-band photometry to select a sample of galaxies in a slice +501 Mpe.," To begin we imagine that we can use spectroscopy or multi-band photometry to select a sample of galaxies in a slice $\pm 50\,h^{-1}$ Mpc."
" At our iducial zcOL or x,c3005 !Mpe. this corresponds ο Asf(l|z)15."," At our fiducial $z\simeq 0.1$, or $\chi_\star\simeq 300\,h^{-1}$ Mpc, this corresponds to $\Delta z/(1+z)\sim 15\%$."
 In this 2D slice we estimate the density using spline kernel interpolation with 4 nearest (in ojection) neighbors., In this 2D slice we estimate the density using spline kernel interpolation with 4 nearest (in projection) neighbors.
 Not surprisinghv. we find that this density is much higher for objects which live in massive idlos than for those which live in smaller halos.," Not surprisingly, we find that this density is much higher for objects which live in massive halos than for those which live in smaller halos."
 As the width of the slice. is increased. the contrast in density between high and low mass halos is reduced. but the trend remains the same.," As the width of the slice is increased the contrast in density between high and low mass halos is reduced, but the trend remains the same."
 Note that our goal is not to optimize the density estimator. but to demonstrate that a useful estimator may be computed even for samples with limited redshift information.," Note that our goal is not to optimize the density estimator, but to demonstrate that a useful estimator may be computed even for samples with limited redshift information."
 Of course. the exact choice of estimator will depend on the data set being considered.," Of course, the exact choice of estimator will depend on the data set being considered."
" To pick a reasonable value of p, we note that halos of 10hTAL. host O(10) galaxies in our models and cover 5(hIMpe)? in projection."," To pick a reasonable value of $\rho_\star$ we note that halos of $10^{15}\,h^{-1}M_\odot$ host $\mathcal{O}(10)$ galaxies in our models and cover $\sim 5\,(h^{-1}{\rm Mpc})^2$ in projection."
 Projected over 505.1 Mpe the background. density is ~(Lh1Mpe)27. so massive: halos are 20 more dense than the mean (p).," Projected over $\pm 50\,h^{-1}$ Mpc the background density is $\sim 0.1\,(h^{-1}{\rm Mpc})^{-2}$, so massive halos are $\sim 20$ more dense than the mean $\bar{\rho}$ )."
" We pick p,=25p as a convenient round number. though our conclusions are not sensitive to the exact choice."," We pick $\rho_\star=25\,\bar{\rho}$ as a convenient round number, though our conclusions are not sensitive to the exact choice."
 Figure 3. shows that this marked correlation function on sub-Alpe scales is dilferent for our two samples reflecting the dillerences in the HOD (Fig. 2))., Figure \ref{fig:mark} shows that this marked correlation function on sub-Mpc scales is different for our two samples reflecting the differences in the HOD (Fig. \ref{fig:hod}) ).
 How cdiscriminatory is this measure?, How discriminatory is this measure?
 For our fiducial volume. Ay?c33 for the two marked correlation functions. compared with X?«1 [or the unweighted correlations.," For our fiducial volume, $\Delta\chi^2\simeq 33$ for the two marked correlation functions, compared with $\Delta\chi^2 < 1$ for the unweighted correlations."
 Since almost all of the difference comes from the lowest 4 cata points. the two models can be stronely discriminated (99% assuming Gaussian errors).," Since almost all of the difference comes from the lowest 4 data points, the two models can be strongly discriminated $>99\%$ assuming Gaussian errors)."
 The distribution of the marks is almost the same in the two samples. and the dillerence in Al(i) remains even if we rescale the marks in one model to match the distribution in the other. showing that the dilference is robust.," The distribution of the marks is almost the same in the two samples, and the difference in $M(r)$ remains even if we rescale the marks in one model to match the distribution in the other, showing that the difference is robust."
 We also note that the relevant measure of error for ALG?) comes from the Monte-Carlo. estimation of the covariance matrix., We also note that the relevant measure of error for $M(r)$ comes from the Monte-Carlo estimation of the covariance matrix.
 Simply scrambling the marks allows us to test for a density dependence of the correlation function citepSheConski05.. which is detected in all of our catalogs at extremely high. significance. but does not tell us how to compare cillerent AZ(r) to each other.," Simply scrambling the marks allows us to test for a density dependence of the correlation function \\citep{SheConSki05}, which is detected in all of our catalogs at extremely high significance, but does not tell us how to compare different $M(r)$ to each other."
 Our initial choice of slice. width. +504tAlpe. was possibly optimistic for surveys at higher redshift.," Our initial choice of slice width, $\pm 50\,h^{-1}$ Mpc, was possibly optimistic for surveys at higher redshift."
 As we increase the width of the slice the density contrast decreases and the significance by which we can cilferentiate the models is also decreased., As we increase the width of the slice the density contrast decreases and the significance by which we can differentiate the models is also decreased.
 For a slice £1251 Mpe in width. ic. the full depth. of our fiducial (250*\Ipe)? survey. using the same mark as above. the two models in Figure 3 diller by AD=19.," For a slice $\pm 125\,h^{-1}$ Mpc in width, i.e. the full depth of our fiducial $(250\,h^{-1}{\rm Mpc})^3$ survey, using the same mark as above, the two models in Figure \ref{fig:mark} differ by $\Delta\chi^2=19$."
" Ht is easily conceivable that a cülferent choice of p, or a higher power of p in the mark could increase the clis¢riminatory power of AJ(r). but this is alreacly reasonably significant. given that only the 4 points with r<Lh !Mpe contribute."," It is easily conceivable that a different choice of $\rho_\star$ or a higher power of $\rho$ in the mark could increase the discriminatory power of $M(r)$, but this is already reasonably significant given that only the 4 points with $r<1\,h^{-1}$ Mpc contribute."
 Lf it proves impossible to select. slices as thin as £1254 *Alpe for some particular sample. it is always possible to jointly analyze samples using Cross-correlation marks where one of the samples can be well isolated. in distance (c.g. red galaxies with strong breaks).," If it proves impossible to select slices as thin as $\pm 125\,h^{-1}$ Mpc for some particular sample, it is always possible to jointly analyze samples using cross-correlation marks where one of the samples can be well isolated in distance (e.g. red galaxies with strong breaks)."
 Such statistics would need to be analvzed on a case-by-case basis., Such statistics would need to be analyzed on a case-by-case basis.
 Although5 the 5galaxy two-point correlation function has proved to be extremely useful in modeling the relationship between 5galaxies and dark matter. it does not exhaust. the information in the data.," Although the galaxy two-point correlation function has proved to be extremely useful in modeling the relationship between galaxies and dark matter, it does not exhaust the information in the data."
 One degeneracy that cannotbe broken by the correlation function. alone is between the LIOD and cosmologyS. - within the context of the correlation function. one is [ree to re-apportion 5galaxies to compensateΙ for cülferences in the halo mass function.," One degeneracy that cannotbe broken by the correlation function alone is between the HOD and cosmology - within the context of the correlation function, one is free to re-apportion galaxies to compensate for differences in the halo mass function."
 The number of such, The number of such
pressure are know: with G. M... and R.. being the eravitational constant. the stellar mass. and stellar racius. respectively.,"pressure are known: with $G$ , $M_\odot$ , and $R_\odot$ being the gravitational constant, the stellar mass, and stellar radius, respectively."
 We then can define the euergy source terni: where Pocos< Pisa racial function of ου., We then can define the energy source term: where $1<\gamma\le \Gamma$ is a radial function of $\gamma_0$.
 It cau be seen that for smaller values of 5 (simaller expausiou of the maeuetic flux tube). the energy source is larger. aud so ds the acceleration of the wind.," It can be seen that for smaller values of $\gamma$ (smaller expansion of the magnetic flux tube), the energy source is larger, and so is the acceleration of the wind."
" At buee distances frou the Sun (~ LOR.) 5ΣΕ and E,>0."," At large distances from the Sun $\sim 10R_\odot$ ), $\gamma\rightarrow\Gamma$, and $E_\gamma \rightarrow 0$."
 The main advantage of this approach over the use of an ad-hoc wine solutions imposed on the maguetic field (the Parker wind solution for example: Parker 1958)) is that here the topology of the wiud depends ou the observed magnetic field structure (c.e..Phillipsetal.1995).," The main advantage of this approach over the use of an ad-hoc wind solutions imposed on the magnetic field (the Parker wind solution for example; \citealt{parker58}) ) is that here the topology of the wind depends on the observed magnetic field structure \citep[e.g.,][]{Phillips95}."
. This model takes iuto account the dependency of the wind structure on the large-scale magnetic topology in the solar/stellar corona aud produces a biauodal wind with regions of slow. denser wind. (which originates frou the opeu/closed field boundary where the expansion factor is large). and regions of fast. less dense wind (which originates from open field regious where the expansion actor is Μπα) (AleComasetal.2007).," This model takes into account the dependency of the wind structure on the large-scale magnetic topology in the solar/stellar corona and produces a bi-modal wind with regions of slow, denser wind, (which originates from the open/closed field boundary where the expansion factor is large), and regions of fast, less dense wind (which originates from open field regions where the expansion factor is small) \citep{McComas07}."
. Iu order to avoid complexity and isolate the planet-CME interaction. we set the stellar magnetic field to ο a dipole with an equatorial field streneth of 2.5 aligned with the rotation axis of the star.," In order to avoid complexity and isolate the planet-CME interaction, we set the stellar magnetic field to be a dipole with an equatorial field strength of 2.5G aligned with the rotation axis of the star."
" Other stellar xuanmeters used Lere are watched to the observed stellar xuanmeters of ΠΟ 1897233. with stellar radius. δν=Πλο stellar mass. M,=0.2.M ... and stellar rotation veriod. O,=11.954 (Schneider1995:Mavoretal.2003)."," Other stellar parameters used here are matched to the observed stellar parameters of HD 189733, with stellar radius, $R_\star=0.76R_\odot$ , stellar mass, $M_\star=0.2M_\odot$ , and stellar rotation period, $\Omega_\star=11.95d$ \citep{exoplanet95,exoplanets03}."
. The ambicut stellar wind is solar-like with terminal speeds ranging between 265Ag band 800hans| (seeCohenetal.20096.201010).," The ambient stellar wind is solar-like with terminal speeds ranging between $265\;km\;s^{-1}$ and $800\;km\;s^{-1}$ \citep[see][]{cohen09b,Cohen10c}."
 Iu order to model the planet. we impose an additional boundary coudition iu the μπαπο domain which is constrained by the planetary surface deusity. temperature. aud magnetic field ina similar miuiner as in Cohenetal.(20000) aud Cohenetal.(20105).," In order to model the planet, we impose an additional boundary condition in the simulation domain which is constrained by the planetary surface density, temperature, and magnetic field in a similar manner as in \cite{cohen09b} and \cite{Cohen10c}."
.. At cach time-step. the coordinates of this boundary (or planetary body) are updated based on the planetary orbit.," At each time-step, the coordinates of this boundary (or planetary body) are updated based on the planetary orbit."
" Cuid cells inside the body are defined as ""ghost cells” or “boundary cells” aud are forced to have the boundary values.", Grid cells inside the body are defined as “ghost cells” or “boundary cells” and are forced to have the boundary values.
" The solution is updated given the particular set of ""boundary cells” at a particular time-step.", The solution is updated given the particular set of “boundary cells” at a particular time-step.
" At cach time step. the planets location is updated. aud boundary cells that no longer overlap the planet are returned to being a ""yegular cell aud their values are updated according to the MIID solution."," At each time step, the planet's location is updated, and boundary cells that no longer overlap the planet are returned to being a “regular cell” and their values are updated according to the MHD solution."
 This process is illustrated in Figure 5.., This process is illustrated in Figure \ref{fig:f1}.
" Tere. the secoud boundary condition for the planet is set at a semi-major axis of &=SAR, aud with a radius of R,=0.2/8,LOR). where Ry is Jupiter's radius."," Here, the second boundary condition for the planet is set at a semi-major axis of $a=8.8R_\star$ and with a radius of $R_p=0.2R_\star\approx 1.5R_J$, where $R_J$ is Jupiter's radius."
" The boundary for the base density and temperature aren,mcouditious=10*en? aud T,-104ke. respectively."," The boundary conditions for the base density and temperature are $n_p=10^7~cm^{-3}$ and $T_p=10^4~k$, respectively."
 Here we do not include the planetary eravitv since several tests have shown that the effect of the simall planetary mass on the solution is negligible., Here we do not include the planetary gravity since several tests have shown that the effect of the small planetary mass on the solution is negligible.
 This is due to the fact that the planet is small compared to the siuulation domain. so that the planetary scale height cannot be captured as also other phivsical features. Huportant ou a planetary scale.," This is due to the fact that the planet is small compared to the simulation domain, so that the planetary scale height cannot be captured as also other physical features, important on a planetary scale."
 The region near the planet is magnetically domiuated. aud so stress balance between the magnetosphere and wind/CME does not depend on density (through thermal pressure. which requires a 120re realistic 1iodel for the planet).," The region near the planet is magnetically dominated, and so stress balance between the magnetosphere and wind/CME does not depend on density (through thermal pressure, which requires a more realistic model for the planet)."
 Therefore. the deusitv eradicut near the plauet is dominated bv the difference between the boundary density value iud the ambient deusitv of the stellar corona. as well as the eradieunt iu planetary magnetic pressure.," Therefore, the density gradient near the planet is dominated by the difference between the boundary density value and the ambient density of the stellar corona, as well as the gradient in planetary magnetic pressure."
" This deusitv eracdient is more moderate than that in realitv. aud it shows. for cxaple. a drop of less than order of magnitude between 75, aud 2f/8,. while iu reality oue should expect a drop 2-3 orders of iaguitude based on the planetary parameters."," This density gradient is more moderate than that in reality, and it shows, for example, a drop of less than order of magnitude between $R_p$ and $2R_p$, while in reality one should expect a drop 2-3 orders of magnitude based on the planetary parameters."
 Planctary rotation is onütted: even though it could be nmipleueuted via the boundary condition for the velocity iuside the secoud body (currently. the boundary condition for the planetary velocity is of the orbital volocitv).," Planetary rotation is omitted; even though it could be implemented via the boundary condition for the velocity inside the second body (currently, the boundary condition for the planetary velocity is of the orbital velocity)."
 We choose to ignore rotation in this simulation in order not to apply more complexity to a nunerical boundary. which already las a delicate stability due to its tiny size.," We choose to ignore rotation in this simulation in order not to apply more complexity to a numerical boundary, which already has a delicate stability due to its tiny size."
 We assume that close-in planets do uot rotate as fast as Jupiter due to their spiudowui by tidal locking processes (Sáuchez-Lavega.2001)... aud that the planetary maguetospheric dvnznuices is Earth-like (dominated bv the solar wind) aud not Jupiter-like (dominated by plauetarv rotation) (INivelsoun&Russell 1995).," We assume that close-in planets do not rotate as fast as Jupiter due to their spindown by tidal locking processes \citep{sanches-lavega04}, and that the planetary magnetospheric dynamics is Earth-like (dominated by the solar wind) and not Jupiter-like (dominated by planetary rotation) \citep{RussellKivelson95}."
". Since our goal is to quantity the planetary magnetic field streugth necessary to shield the planet from CME events, we study two cases;"," Since our goal is to quantify the planetary magnetic field strength necessary to shield the planet from CME events, we study two cases."
 One with a weak equatorial dipole field streneth of 0.5G (Case A hereafter) aud one with a stronger equatorial dipole field streneth of 1C (Case B hereafter)., One with a weak equatorial dipole field strength of $0.5~G$ (Case A hereafter) and one with a stronger equatorial dipole field strength of $1~G$ (Case B hereafter).
 Jupiters equatorial field is about 1.3G., Jupiter's equatorial field is about $4.3~G$.
 We start the simulation bv lettiue the win solution relax to a steady-state with a stationary (tically locked) plauet aud then turn ou the planetary orbita motion., We start the simulation by letting the wind solution relax to a steady-state with a stationary (tidally locked) planet and then turn on the planetary orbital motion.
 The orbital motion is obtained by updating the coordinates of the planetary boundary at cach time-step., The orbital motion is obtained by updating the coordinates of the planetary boundary at each time-step.
 This technique aud its implications for the cvnamics of stellar coronae harboriug close-in planets is describe: in Cohenoetal(2011)., This technique and its implications for the dynamics of stellar coronae harboring close-in planets is described in \cite{Cohen11}.
. We allow the simulation to evolve for half an orbit so that it does not inehide auv verturbations generated by the initiation of the orbita notion., We allow the simulation to evolve for half an orbit so that it does not include any perturbations generated by the initiation of the orbital motion.
 We then use this solution as our initial condition and launch the CALE as described below refsec:CAIE))., We then use this solution as our initial condition and launch the CME as described below \\ref{sec:CME}) ).
 Figure Ll) shows the distribution of the nuuber density aud temperature at this pre-cruption state over the equatorial plain of the simulation domain., Figure \ref{fig:f2} shows the distribution of the number density and temperature at this pre-eruption state over the equatorial plain of the simulation domain.
 The relatively Ligh temperature of ~10°A in the jlanctary tail (comparing to the boundary value of LO! WW) is due to the fact that it containshotcoronal asma that is beimg trapped inside the planctarynaenetosplere as it is sweeping through the corona., The relatively high temperature of $\sim 10^5\;K$ in the planetary tail (comparing to the boundary value of $10^4\;K$ ) is due to the fact that it containshotcoronal plasma that is being trapped inside the planetarymagnetosphere as it is sweeping through the corona.
 We initiate the CME by superimposing au unstable. semi-circular flux rope based on the analytical model by Titov&Déómouliu(1999). ou top of the ambicut “initial” solution described above," We initiate the CME by superimposing an unstable, semi-circular flux rope based on the analytical model by \cite{titov99} on top of the ambient “initial” solution described above"
" (Lada&Lada2003: Melkee&Ostriker2007)). 200030%. ofAL)=dN/dAL. c(AJ)xAL’. 10!AJ, 1093/. 2mLz (Rosolowsky2005:Blitzetal.2007:Fukui2008).."," \citealt{lada03}; \citealt{mckee07b}) $20-30\%$ $\psi(M) \equiv dN/dM$ $\psi(M) \propto M^{\beta}$ $10^4M_{\odot}$ $10^6M_{\odot}$ $\beta \approx -1.7$ \citep{rosolowsky05b, blitz07a, fukui08a}."
" or higher-density tracers such as CLO, PCO. and thermal dust emission (Reid&Wilson2006:AIunozetal.2007:Wong 2008).."," or higher-density tracers such as $^{18}$ O, $^{13}$ CO, and thermal dust emission \citep{reid06b, munoz07a, wong08a}."
 Youne star clusters rave sc—2.0 (Ehuegreen&Efremov1997:Mclxeeetal.2009:Clhiaudar 2010).," Young star clusters have $\beta \approx -2.0$ \citep{elmegreen97a, mckee97, zhang99b, dowell08a, 
fall09a, chandar09a}."
".. The similar exponeuts or clouds aud clusters indicate that the efficiency of star ormation and probability of disuption are at most weal ""uctious of mass.", The similar exponents for clouds and clusters indicate that the efficiency of star formation and probability of disruption are at most weak functions of mass.
 This couclusion is reiuforced. by the aet that 3 is the same for 101LOS s3-0ld clusters as if is for 10—10* yr-olel clusters (Zhang&Fall1999:etal.2009:Chandarct 2010).," This conclusion is reinforced by the fact that $\beta$ is the same for $10^7-10^8$ yr-old clusters as it is for $10^6-10^7$ yr-old clusters \citep{zhang99b, fall09a, 
chandar09a}."
. These enipirical results may at first seen puzzliug., These empirical results may at first seem puzzling.
 Low-nass protoclusters have lower binding cnerev pcr tit mass and should therefore be easier to disrupt tlauhighanass protoclusters., Low-mass protoclusters have lower binding energy per unit mass and should therefore be easier to disrupt thanhigh-mass protoclusters.
 Tudecd. several authors have proposed that feedback would cause a bend in the mass function of voung clusters at AL~107AZ... motivated in part by the well-known turnover m the mass function of old globular clusters (sroupa&Boily2002:Batiu-eardtetal.2008:Paxiueutier 2008)..," Indeed, several authors have proposed that feedback would cause a bend in the mass function of young clusters at $M \sim 10^5M_{\odot}$, motivated in part by the well-known turnover in the mass function of old globular clusters \citep{kroupa02a, baumgardt08a, parmentier08a}."
 For voung clusters. such a feature is not observed (as noted above). while for globular clusters. it arises from almost any initial conditions as a cousequenuce of stellar escape driven bv two-body relaxation over ~10/9 vy (Fall&Zhaug2001:MeLbaughliu&Fall2008.andreferences therein)...," For young clusters, such a feature is not observed (as noted above), while for globular clusters, it arises from almost any initial conditions as a consequence of stellar escape driven by two-body relaxation over $\sim 10^{10}$ yr \citep[and references therein]{fall01a, mclaughlin08a}."
 Nevertheless. we are left with an müportaut question: What are the physical reasous for the observed similarity of the mass fictions of iiolecular clouds aud voung star ‘listers?," Nevertheless, we are left with an important question: What are the physical reasons for the observed similarity of the mass functions of molecular clouds and young star clusters?"
 The goal of this Letter is to answer this question., The goal of this Letter is to answer this question.
 Iu Section 2.. we derive some general relations. betweenthemass functionsof cloudsaudclusters.," In Section \ref{energymomentum}, we derive some general relations betweenthemass functionsof cloudsandclusters."
 In Section 3.. we review a varicty of specific feedback processes aud estimate the star formation efficiency for radiation," In Section \ref{sec:efficiency}, , we review a variety of specific feedback processes and estimate the star formation efficiency for radiation"
We model spin-down of a pulsar under the combined action of magnetic dipole radiation and propeller spin-down torques as adopting the model of Menou. Perna and IHernequist (2001b) with their notation.,"We model spin-down of a pulsar under the combined action of magnetic dipole radiation and propeller spin-down torques as adopting the model of Menou, Perna and Hernquist (2001b) with their notation."
 The neutron star will continue to act as a radio pulsar as long as (he Lall-back disk does nol protrude into the lisht exlinder., The neutron star will continue to act as a radio pulsar as long as the fall-back disk does not protrude into the light cylinder.
 If the disk were detached from the light evlinder. it would not exert anv torque on the neutron star ancl its magnetosphere.," If the disk were detached from the light cylinder, it would not exert any torque on the neutron star and its magnetosphere."
 Assuming that the disk is altached to the light evlinder. (he torque can be estimated as: where M is the mass inflow rate interacting with the light evlinder and being ejected [rom the disk: rj.=e/Q is the light cylinder radius. and ΤΟ is the specilic angular momentum extracted [rom the pulsar magnetosphere. since the Neplerian rotation rate (ή) in the disk is small compared to the rotation rate O of the neutron star and its magnetosphere.," Assuming that the disk is attached to the light cylinder, the torque can be estimated as: where $\dot{M}$ is the mass inflow rate interacting with the light cylinder and being ejected from the disk; $r_{lc}=c/\Omega$ is the light cylinder radius, and $r_{lc}^2\Omega$ is the specific angular momentum extracted from the pulsar magnetosphere, since the Keplerian rotation rate $\Omega_K(r_{lc})$ in the disk is small compared to the rotation rate $\Omega$ of the neutron star and its magnetosphere."
" The parameter 5=ολο—2x10M, erg/s is (he rate of energy loss of (he neutron star due to the propeller torque: qi is the mass inflow rate in units of 10! em/s. The rate of enerev loss due to magnetic dipole radiation is given by lere D. is the component of the dipole magnetic field at the neutron star surface in the direction perpendicular to the rotation axis aud Rois the neutron star radius."," The parameter $\gamma=2\dot{M} c^2 = 2 \times 10^{31} \dot{M}_{10}\,$ erg/s is the rate of energy loss of the neutron star due to the propeller torque; $\dot{M}_{10}$ is the mass inflow rate in units of $10^{10}$ gm/s. The rate of energy loss due to magnetic dipole radiation is given by Here $B_\bot$ is the component of the dipole magnetic field at the neutron star surface in the direction perpendicular to the rotation axis and R is the neutron star radius."
 This defines j—6617x107οR.," This defines $\beta=6.17 \times 10^{27} {B_{\bot,12}}^2 {R_6}^6$."
 Menou. Perna and Hernquist (2001b) give the solution of Eq.(1) for constant AM as where Q;. the initial rotation rate of the pulsar. is always large enough to justify the second equation.," Menou, Perna and Hernquist (2001b) give the solution of Eq.(1) for constant $\dot{M}$ as where $\Omega_i$, the initial rotation rate of the pulsar, is always large enough to justify the second equation."
 The timescale 7 is 7=1/2(53)/?4.5x1072444Doy| ves.," The timescale $\tau$ is $\tau=I/2(\gamma\beta)^{1/2}= 4.5 \times 10^7 
{\dot{M}_{10}}^{-1/2}{B_{\bot,12}}^{-1}$ yrs."
 In the P— diagram pulsars will follow (racks given by Equation (1).," In the $P-\dot{P}$ diagram pulsars will follow tracks given by Equation (1),"
near-IR and the observed visibilities; this model can therefore be ruled out.,near-IR and the observed visibilities; this model can therefore be ruled out.
 We have performed a detailed radiative transfer modeling of the observed SED and visibilities with the RADMC code., We have performed a detailed radiative transfer modeling of the observed SED and visibilities with the RADMC code.
 It suggests a scenario of a geometrically flat but optically thick circumstellar dust disk., It suggests a scenario of a geometrically flat but optically thick circumstellar dust disk.
" We also used the fitting tool by ?,, which leads to similar disk models and therefore confirms the results obtained with RADMC, whereas envelope models found within the Robitaille grid did not provide similarly good fits to SED and visibilities simultaneously."," We also used the fitting tool by \citet{Robitaille}, which leads to similar disk models and therefore confirms the results obtained with RADMC, whereas envelope models found within the Robitaille grid did not provide similarly good fits to SED and visibilities simultaneously."
" Although there is no observational evidence, the possible presence of a large (r5400 AU) spherical envelope around the disk cannot be excluded by the data."," Although there is no observational evidence, the possible presence of a large $r \geq 400$ AU) spherical envelope around the disk cannot be excluded by the data."
 The mass of the circumstellar disk is about 0.1Μο.," The mass of the circumstellar disk is about $0.1\,M_\odot$."
 The data suggest an inclination angle of ~30° and a position angle of ~40° for the orientation of the disk., The data suggest an inclination angle of $\sim 30\degr$ and a position angle of $\sim 40\degr$ for the orientation of the disk.
 These values are consistent with an overall geometrical model based on the jet-like feature seen to the north-east of IRS 1., These values are consistent with an overall geometrical model based on the jet-like feature seen to the north-east of IRS 1.
" Comparing NGC 2264 IRS 1 to other YSOs, the size of its MIR emitting region seems to be typical of its luminosity and mass."," Comparing NGC 2264 IRS 1 to other YSOs, the size of its MIR emitting region seems to be typical of its luminosity and mass."
" However, the large scatter of sizes in this range of luminosities and masses points towards a wide variety of (disk) morphologies among these objects."," However, the large scatter of sizes in this range of luminosities and masses points towards a wide variety of (disk) morphologies among these objects."
" More observational research, such as future MIR interferometric observations, will provide tighter constraints on the circumstellar material around this interesting source and therefore will help us to clarify our understanding of the disks around high-mass YSOs."," More observational research, such as future MIR interferometric observations, will provide tighter constraints on the circumstellar material around this interesting source and therefore will help us to clarify our understanding of the disks around high-mass YSOs."
Cepheid P-L relation suggests the distance modulus of (m—M)o=24.38+0.05 to M31 when they performed the results of theHST Distance Scale Key Project to measure the Hubble constant.,Cepheid P–L relation suggests the distance modulus of $(m-M)_{0}=24.38\pm0.05$ to M31 when they performed the results of the Distance Scale Key Project to measure the Hubble constant.
 Durrelletal.(2001) determined the distance modulus of (m—M)o=24.47+0.12 to M31 from the luminosity of the RGB tip of over 2000 RGB halo stars in a halo field located about 20 kpc from the M31 nucleus along the southeast minor axis., \citet{Durrell01} determined the distance modulus of $(m-M)_{0}=24.47\pm0.12$ to M31 from the luminosity of the RGB tip of over 2000 RGB halo stars in a halo field located about 20 kpc from the M31 nucleus along the southeast minor axis.
 Joshietal. have obtained R— and I—band observations of a (2003)13’x region in the disk of M31 and derived the Cepheid period-luminosity distance modulus to be (m—M)o=24.49+ 0.11., \citet{Joshi03} have obtained $R-$ and $I-$ band observations of a $13'\times13'$ region in the disk of M31 and derived the Cepheid period–luminosity distance modulus to be $(m-M)_{0}=24.49\pm0.11$ .
 Brownetal. determined the distance modulus of (m—M)o(2004b)=24.5-Ε0.1 to M31 based on brightness of 55 RR Lyrae stars detected on the HST/ACS images of ~ 84 hr (250 exposures over 41 ," \citet{brown04b} determined the distance modulus of $(m-M)_{0}=24.5\pm
0.1$ to M31 based on brightness of 55 RR Lyrae stars detected on the /ACS images of $\sim$ 84 hr (250 exposures over 41 days)."
McConnachieetal. derived the distance days).modulus to M31 to be (m—(2005)M)o=24.47+0.07 based on the method of the tip of the RGB observed using the Isaac Newton Telescope Wide Field Camera (INT WEC)., \citet{McConnachie05} derived the distance modulus to M31 to be $(m-M)_{0}=24.47\pm0.07$ based on the method of the tip of the RGB observed using the Isaac Newton Telescope Wide Field Camera (INT WFC).
 Ribasetal.(2005) derived the distance modulus of M31 as (m—M)o=24.44+0.12 from an eclipsing binary., \citet{ribas05} derived the distance modulus of M31 as $(m-M)_{0}=24.44\pm0.12$ from an eclipsing binary.
" Very recently, Sarajedinietal.(2009) presented theHST observations taken with the ACS WFC of two fields near M32 located 4—6 kpc from the center of M31, and identified 752 RR variables with excellent photometric and temporal completeness."," Very recently, \citet{sarajedini09}
 presented the observations taken with the ACS WFC of two fields near M32 located $4-6$ kpc from the center of M31, and identified 752 RR variables with excellent photometric and temporal completeness."
" Based on this large sample of M31 RR Lyrae variables, and using a relation between RR Lyrae luminosity and metallicity along with a reddening value of E(B—V)=0.08+0.03, they derived the distance modulus of (m—Μο=24.46+0.11 to M31."," Based on this large sample of M31 RR Lyrae variables, and using a relation between RR Lyrae luminosity and metallicity along with a reddening value of $E(B-V)=0.08\pm0.03$, they derived the distance modulus of $(m-M)_0=24.46\pm0.11$ to M31."
" In order to see clearly, we list these determinations of M31 distance moduli in Table 1."," In order to see clearly, we list these determinations of M31 distance moduli in Table 1."
 It is evident that our determination is in good agreement with the previous determinations., It is evident that our determination is in good agreement with the previous determinations.
" In this paper, we re-determined the age of the M31 GC"," In this paper, we re-determined the age of the M31 GC"
rest-frame UV continuum. and also galaxies fainter than our selection limit will also contribute to the integrated UV light. density.,"rest-frame UV continuum, and also galaxies fainter than our selection limit will also contribute to the integrated UV light density."
. Lf D.instead. we integrate. down to AMI!pÉ=13 (corresponding to 0.01AZ.ve +) and the total Xstar formation rate density is 0.0081AZ.ve+Alpe7.," If instead we integrate down to $M^{UV}_{\rm 1600\,\AA }=-13$ (corresponding to $0.01\,M_{\odot}\,{\rm yr^{-1}}$ ) and the total star formation rate density is $0.0081\,M_{\odot}\,{\mathrm{yr}}^{-1}\,{\mathrm{Mpc}}^{-3}$."
 These star formation rate densities are a [actor of ~10lower than at το38 4. and even a factor of z23.5 below that at z6 (Bunker et 22004: Bouwens ct 22006).," These star formation rate densities are a factor of $\sim 10$ than at $z\sim 3-4$ , and even a factor of $\approx 3-5$ below that at $z\approx 6$ (Bunker et 2004; Bouwens et 2006)."
 The ionizing UV. photons produced. by the most massive (OB) stars might be critical in reionization and keeping the Universe ionized at zπο., The ionizing UV photons produced by the most massive (OB) stars might be critical in reionization and keeping the Universe ionized at $z\approx 6-11$.
" However. work at 2σε6 has shown that under standard assumptions of the LME. escape fraction and. clumping of the gas. the observed. population of Lyman break galaxies produce insullicient Lux down to ABs28.5 mmag (Bunker et 22004). and the ""photon drought is even more severe at zcY (Wilkins et 22010b)."," However, work at $z\approx 6$ has shown that under standard assumptions of the IMF, escape fraction and clumping of the gas, the observed population of Lyman break galaxies produce insufficient flux down to $AB\approx 28.5$ mag (Bunker et 2004), and the “photon drought"" is even more severe at $z\approx 7$ (Wilkins et 2010b)."
 We now compare our measured UM. luminosity. density at 2&δ9 (quoted. above as a corresponding star formation rate densitv) with that required. to. ionize the Universe at this recshift., We now compare our measured UV luminosity density at $z\approx 8-9$ (quoted above as a corresponding star formation rate density) with that required to ionize the Universe at this redshift.
 Madau. Llaardt Rees (1999) eive the density of star formation required. for reionization (assuming the same Salpeter IME as used in this paper): We have updated equation 27 of Macau. Laardt Rees (1999) [or a more recent concordance cosmology estimate of the barvon density Crom Larson et (2010). OyΤρι= 0.022622.," Madau, Haardt Rees (1999) give the density of star formation required for reionization (assuming the same Salpeter IMF as used in this paper): We have updated equation 27 of Madau, Haardt Rees (1999) for a more recent concordance cosmology estimate of the baryon density from Larson et (2010), $\Omega_b\,h_{100}^2=0.022622$ ."
 The reionization requirement at z8.6 is a factor of 2.5 times higher than that at +=6. as the number of photons needed rises as (1]z)5," The reionization requirement at $z\approx 8.6$ is a factor of 2.5 times higher than that at $z\approx 6$, as the number of photons needed rises as $(1+z)^3$."
 In the above equation. Co is the clumping factor. of neutral hydrogen. €=pj)(pma)7.," In the above equation, $C$ is the clumping factor of neutral hydrogen, $C=\left< \rho^{2}_{\mathrm{HI}}\right> \left< \rho_{\mathrm{HI}}\right> ^{-2}$."
 Early simulations suggested CC30 (CGnedin Ostriker LOOT). but more recent work including the ellects of reheating implies a lower concentration factor of Cox5 (Pawlik et 22009).," Early simulations suggested $C\approx 30$ (Gnedin Ostriker 1997), but more recent work including the effects of reheating implies a lower concentration factor of $C\approx 5$ (Pawlik et 2009)."
 The escape fraction of ionizing whotons Cfi) for high-redshift galaxies is highly. uncertain (c.g.. Steidel. Pettini Acdelberger 2001. Shapley et 22006). and it is possible hat the escape fraction of ionising photons may be linked o the escape fraction of Lyman-a photons (Stark ct 22010). which may mean that high. escape fractions could oe tested through future line emission line searches with spectroscopy ancl narrow-bancl imaging.," The escape fraction of ionizing photons $f_{\mathrm{esc}}$ ) for high-redshift galaxies is highly uncertain (e.g., Steidel, Pettini Adelberger 2001, Shapley et 2006), and it is possible that the escape fraction of ionising photons may be linked to the escape fraction of $\alpha$ photons (Stark et 2010), which may mean that high escape fractions could be tested through future line emission line searches with spectroscopy and narrow-band imaging."
 Even if we take he upper limit of fice=1 (no absorption by 1) and a very low clumping factor. the required total star formation rate density for reionization is 0.012AZ.ve1Alpe *.," Even if we take the upper limit of $f_{\mathrm{esc}}=1$ (no absorption by ) and a very low clumping factor, the required total star formation rate density for reionization is $0.012\,M_{\odot}\,{\mathrm{yr}}^{-1}\,{\mathrm{Mpc}}^{
-3}$ ."
 This is a [actor of ~5 higher than our measured star formation density at z89 from Y-drop galaxies brighter than Alpe=18.5 (our approximate limit)., This is a factor of $\sim 5$ higher than our measured star formation density at $z\approx 8-9$ from $Y$ -drop galaxies brighter than $M_{UV}=-18.5$ (our approximate limit).
" As shown in Table and Figure δι, the required UV Luminosity density can only just be achieved (if fi= 1) by integrating down to Mi:=13 (Le. extrapolating the Schechter function to z100 times fainter than our observed limit) and then only for a steeper faint end slope ofa=1.9 rather than aΞ1.3."," As shown in Table \ref{tab:uvld} and Figure \ref{fig:sfrd_Y}, the required UV luminosity density can only just be achieved (if $f_{\mathrm{esc}}=1$ ) by integrating down to $M_{UV}=-13$ (i.e., extrapolating the Schechter function to $\approx 100$ times fainter than our observed limit) and then only for a steeper faint end slope of $\alpha=-1.9$ rather than $\alpha=-1.7$."
 Aclopting a less unrealistic value of fi;=0.7 (which is still high compared. with observed: values at lower redshift) the required. total star. formation rate density for reionization would be 0.017AZ.ve1Alpe* then the Y«drop population can onlv provide sullicient ionizing photons if the faint end slope is very. steep (a1.9) and the Schechter function is integrated. down below Adpy= δ (corresponding to a star formation rate of only 107A7.vr.1).," Adopting a less unrealistic value of $f_{\mathrm{esc}}=0.7$ (which is still high compared with observed values at lower redshift) the required total star formation rate density for reionization would be $0.017\,M_{\odot}\,{\mathrm{yr}}^{-1}\,{\mathrm{Mpc}}^{
-3}$, then the $Y$ -drop population can only provide sufficient ionizing photons if the faint end slope is very steep $\alpha\le-1.9$ ) and the Schechter function is integrated down below $M_{UV}=-8$ (corresponding to a star formation rate of only $10^{-4}\,M_{\odot}\,{\mathrm{yr}}^{-1}$ )."
 We note that recent theoretical papers indicate that the reionization process itself may have been “photon-starved™ (¢.g.. Bolton lHlaehnelt 2007). consistent with the extrapolation of our observational constraints.," We note that recent theoretical papers indicate that the reionization process itself may have been “photon-starved"" (e.g., Bolton Haehnelt 2007), consistent with the extrapolation of our observational constraints."
 Llowever. the assumption of a solar metallicity Salpeter IME may be flawed: the colours of 2~6 7-band drop-outs are very blue (Stanway. MeMahlon Bunker 2005). with p<2. and the recent WECS J- and Z-band images show that the zzτ z-drops also have blue colours on average (Bunker et 22010: Bouwens et 22010b: Wilkins et 22010c).," However, the assumption of a solar metallicity Salpeter IMF may be flawed: the colours of $z\sim 6$ $i'$ -band drop-outs are very blue (Stanway, McMahon Bunker 2005), with $\beta<-2$, and the recent WFC3 $J$ - and $H$ -band images show that the $z\approx 7$ $z'$ -drops also have blue colours on average (Bunker et 2010; Bouwens et 2010b; Wilkins et 2010c)."
 Continuous star formation with a Salpeter IME produces a UV. spectral slope of 3=2 if there is no dust. redcdening., Continuous star formation with a Salpeter IMF produces a UV spectral slope of $\beta\approx -2$ if there is no dust reddening.
 The [act that we observe even more blue slopes than this (43<2) could be explained through low metallicity. or a top-heavy EME: which can produce between 53 and LO times as many ionizing photons for the same UUV luminosity (Schaerer 2003. see also Stiavelli. Fall Panagia 2004).," The fact that we observe even more blue slopes than this $\beta<-2$ ) could be explained through low metallicity, or a top-heavy IMF, which can produce between 3 and 10 times as many ionizing photons for the same UV luminosity (Schaerer 2003 – see also Stiavelli, Fall Panagia 2004)."
 Alternatively. we may be seeing galaxies at the onset of star formation. or with a rising star formation rate (Verma et 22007). which would. also load us to unclerestimate the true star formation rate from the rest-UV Luminosity.," Alternatively, we may be seeing galaxies at the onset of star formation, or with a rising star formation rate (Verma et 2007), which would also lead us to underestimate the true star formation rate from the rest-UV luminosity."
 We explore the implications of the blue UV spectral slopes in z26 galaxies in a forthcoming paper (Wilkins et 22010c)., We explore the implications of the blue UV spectral slopes in $z\ge 6$ galaxies in a forthcoming paper (Wilkins et 2010c).
 In this paper we have presented. a search forgalaxies at 7.5«z10 using the latest4757 WECS near-infrared data. based on the Lyman-break technique.," In this paper we have presented a search forgalaxies at $7.5<z<10$ using the latest WFC3 near-infrared data, based on the Lyman-break technique."
 Searching for galaxies which have Large (YJ) colours. (Y-drops) on account of the Lyman-alpha forest. absorption. anc with (//——44) colours inconsistent with being Iow-redshift contaminants. we identify =20 candidates at redshift οzzS.9 over an area of =50 square areminutes.," Searching for galaxies which have large $(Y-J)$ colours $Y$ -drops) on account of the Lyman-alpha forest absorption, and with $(J-H)$ colours inconsistent with being low-redshift contaminants, we identify $\approx 20$ candidates at redshift $z \approx 8 - 9$ over an area of $\approx 50$ square arcminutes."
 Our deepest field (the ICDL. Covering 4.2aarcmin?) reaches Jig=28.5 at 660. while the wide-area ERS data (comprising LO WIC3 pointings 'overing 3Taarcmiün?) reaches Jig=27.2.," Our deepest field (the HUDF, covering $^2$ ) reaches $J_{AB}=28.5$ at $6\,\sigma$, while the wide-area ERS data (comprising 10 WFC3 pointings covering $^2$ ) reaches $J_{AB}=27.2$."
 The surface densities of candidates as a function of limiting magnitude appear broadly consistent between our + fields. although jese all lic within I0aaremin.," The surface densities of candidates as a function of limiting magnitude appear broadly consistent between our 4 fields, although these all lie within arcmin."
 Previous searches for Y-drops with WECS have focussed only on the HUDE. and our larger survey has trebelled the number of robust. Y-rop candidates. as well as providing a number of brighter Y-drops (with Jig%21.0 rather than J4g228.0 as in the IIUDE).These brighter sources may be more amenable tospectroscopic follow-up.," Previous searches for $Y$ -drops with WFC3 have focussed only on the HUDF, and our larger survey has trebelled the number of robust $Y$ -drop candidates, as well as providing a number of brighter $Y$ -drops (with $J_{AB}\approx 27.0$ rather than $J_{AB}>28.0$ as in the HUDF).These brighter sources may be more amenable tospectroscopic follow-up."
 For the first time. we have a sullicient number of DEN9 galaxies to fit ó* and M assuming a Schechter luminosity function (previous estimates had. to fix one of these parameters).," For the first time, we have a sufficient number of $z\approx 8-9$ galaxies to fit $\phi^*$ and $M^*$ assuming a Schechter luminosity function (previous estimates had to fix one of these parameters)."
 We confirm that there is large evolution from z= 3. particularly in the bright end of the luminosity function. in the sense that there are lar fewer UV-bright galaxies al 28S9 than in the more recentpast.," We confirm that there is large evolution from $z=3$ , particularly in the bright end of the luminosity function, in the sense that there are far fewer UV-bright galaxies at $z\approx 8-9$ than in the more recentpast."
The cussion properties of the neutron stars surface have beeu first analyzed by Lenzen&Tritmper(1978). and iu some more detail by Brinkmann (1980. hereafter D80).,"The emission properties of the neutron stars surface have been first analyzed by \cite{trule78} and in some more detail by Brinkmann (1980, hereafter B80)."
 Both of these works were aimed to X-ray pulsars. where the surface telperature is a few keVs. aud treated the medium inside the star as a cold electron plasima. ucelecting all possible οπουτς due to electron degeneracy aud ion lattice (primarily through electron-phonon interactions).," Both of these works were aimed to X-ray pulsars, where the surface temperature is a few keV's, and treated the medium inside the star as a cold electron plasma, neglecting all possible effects due to electron degeneracy and ion lattice (primarily through electron-phonon interactions)."
 The (coustaut) damping frequency which appears in DS8O calculations is mainly used to smear the resonance at the evclotrou frequency., The (constant) damping frequency which appears in B80 calculations is mainly used to smear the resonance at the cyclotron frequency.
 Moreover. birefringence in the maguetized vacuum outside the star was not accounted for.," Moreover, birefringence in the magnetized vacuum outside the star was not accounted for."
 In this section we derive the NS surface enissivitv following an approach simular to that discussed in DSO., In this section we derive the NS surface emissivity following an approach similar to that discussed in B80.
 To better illustrate the importance of clectrou-phonon interactions. we first consider a pure cold electron plasma. repeating Druknmonus calculation for the parameter values appropriate to cold isolated NSs rofeold)).," To better illustrate the importance of electron-phonon interactions, we first consider a pure cold electron plasma, repeating Brinkmann's calculation for the parameter values appropriate to cold isolated NSs \\ref{cold}) )."
 A complete treatment which includes the polarization properties of magnetized vacuum is presented in Appendix A.., A complete treatment which includes the polarization properties of magnetized vacuum is presented in Appendix \ref{app1}.
 Since. as we show there. this more general approach is quite cumbersome and ouly eives tiny differeuces with respect to the simpler oue based on unpolarized radiation. the latter is used below.," Since, as we show there, this more general approach is quite cumbersome and only gives tiny differences with respect to the simpler one based on unpolarized radiation, the latter is used below."
 Iu we analyze the more realistic case in which the damping of electromagnetic waves produced by the presence of the ion lattice is iucluded., In \\ref{damp} we analyze the more realistic case in which the damping of electromagnetic waves produced by the presence of the ion lattice is included.
 We start considering the medimm inside the star as a cold electron plasma and ucelect the damping of free electrous due to collisions., We start considering the medium inside the star as a cold electron plasma and neglect the damping of free electrons due to collisions.
 We introduce a cartesian frame as in D80 (see his figures 1 and 2) with the z-axis parallel to the surface normal., We introduce a cartesian frame as in B80 (see his figures 1 and 2) with the $z$ -axis parallel to the surface normal.
 The direction of the incident wave vector k is specified by the angele of incidence / aud the azimuth ο., The direction of the incident wave vector $\mathbf{k}$ is specified by the angle of incidence $i$ and the azimuth $\beta$.
 The magneticfield direction b=B/B is at an angle à with respect to the z-axis and b lies in the «12 plane., The magneticfield direction ${ \bf b} \equiv \mathbf{B} /B$ is at an angle $\alpha$ with respect to the $z$ -axis and $\mathbf{b}$ lies in the $x-z$ plane.
" Given a star surface clement 4A=22R?sinddé at maguctic co-latitude 0. we first compute the total reflectivity p, of the surface for incident unpolarized radiation."," Given a star surface element $dA=2\pi R^2\sin\theta\, d\theta$ at magnetic co-latitude $\theta$, we first compute the total reflectivity $\rho_\omega$ of the surface for incident unpolarized radiation."
" Then. since the absorption cocfiicicut is simply à,=1py. Iirchlioff's ww yields the chussivity 4,=a4I,C where T is the temperature of the emitting clement."," Then, since the absorption coefficient is simply $\alpha_\omega=1-\rho_\omega$, Kirchhoff's law yields the emissivity $j_\omega = \alpha_\omega
B_\omega(T)$, where $T$ is the temperature of the emitting element."
" In general. e, depends on the direction of the refracted rav (see below)."," In general, $\rho_\omega$ depends on the direction of the refracted ray (see below)."
"P). Therefore. the monochromatic flux fi. cmitted by the surface clement τις be computed by iutegratiug over all incident directious. The flux euütted by the eutire surface is given by At the surface. an incident clectromagnetic wave. described by its electric field E aud wave vector k. is partly reflected (E"". Kk"") aud partly refracted."," Therefore, the monochromatic flux $f_\omega$ emitted by the surface element must be computed by integrating over all incident directions, The flux emitted by the entire surface is given by At the surface, an incident electromagnetic wave, described by its electric field ${\bf E}$ and wave vector ${\bf k}$, is partly reflected ${\bf E}'', {\bf k}''$ ) and partly refracted."
 Due to the birefringence of the τοαι. the refracted wave is the sum of au ordinary (EA. and an extraordinay (E. k5) mode.," Due to the birefringence of the medium, the refracted wave is the sum of an ordinary ${\bf E}'_1, {\bf k}'_1$ ) and an extraordinary ${\bf
E}'_2, {\bf k}'_2$ ) mode."
 In order to compute the reflectivity. we necd to solve the dispersion relation and to KA)compute the refractive iudex » for the two modes of propagation.," In order to compute the reflectivity, we need to solve the dispersion relation and to compute the refractive index $n$ for the two modes of propagation."
 In our frame. the diclectiic tensor for a cold electron plasma is given by with By introducingB the Maxwell tensor Ajj=MkjUc|(erP/e7)e;g. where A aro the cartesianB components of+ k’ aud j=ME. the dispersion relation is obtaiued by imposing A;)|= 0.," In our frame, the dielectric tensor for a cold electron plasma is given by with By introducing the Maxwell tensor $\lambda_{ij}= k'_ik'_j-
k^{\prime 2}\delta_{ij}+(\omega^2/c^2)\epsilon_{ij}$, where $k'_i$ are the cartesian components of ${\bf k'}$ and $k^{\prime 2}\equiv k'_ik'_i$, the dispersion relation is obtained by imposing $\vert\lambda_{ij}\vert=0$ ."
 For our purposes it is convenient to write the resulting expression in terms of auele of incidence 7. aud the (complex) refractive index à2h’e fw.," For our purposes it is convenient to write the resulting expression in terms of angle of incidence $i$, and the (complex) refractive index $n=k' c/\omega$ ."
 By using au expression formally analogous to Suell's lw (0=sin//sinO (where now Ο is a complex quantity which replaces the anele of refraction while His real: see e.g. Marion 1965)). it is," By using an expression formally analogous to Snell's law $n = \sin i
/\sin \Theta$ (where now $\Theta$ is a complex quantity which replaces the angle of refraction while$i$ is real; see e.g. \citealt{marion}) ), it is"
"In addition to comparing the central entropy of the local and distant clusters, we also investigated a correlation between K20 and cs.","In addition to comparing the central entropy of the local and distant clusters, we also investigated a correlation between K20 and $c_{SB}$."
 The result is presented in Fig., The result is presented in Fig.
 11 and shows a strong anti-correlation between these two quantities., 11 and shows a strong anti-correlation between these two quantities.
 A Spearman rank test confirms this relation with a coefficient p=—0.84., A Spearman rank test confirms this relation with a coefficient $\rho$ $-$ 0.84.
" As mentioned earlier, we used the global cluster temperature to compute K20, which may introducea bias in the result, in particular for the local clusters that are well resolved down to a few kpc."," As mentioned earlier, we used the global cluster temperature to compute K20, which may introduce a bias in the result, in particular for the local clusters that are well resolved down to a few kpc."
" Since all local clusters were independently analyzed by Cavagnolo et al.,"," Since all local clusters were independently analyzed by Cavagnolo et al.,"
 we used the cluster temperatures measured in the innermost bin (corresponding to about 10-20 kpc) and compared the histograms of K20 using both the central and global temperatures., we used the cluster temperatures measured in the innermost bin (corresponding to about 10-20 kpc) and compared the histograms of K20 using both the central and global temperatures.
 The average difference between the two distributions is 15.3 keV/cm? with a standard deviation of 32.0 keV/cm?., The average difference between the two distributions is 15.3 $^{2}$ with a standard deviation of 32.0 $^{2}$.
" These distributions differ mainly for the nine clusters with K20 «100 keV/cm?, in which K20 (T giopal) 15 On average two times higher than K20 (Tore)."," These distributions differ mainly for the nine clusters with K20 $<$ 100 $^{2}$, in which K20 $T_{global}$ ) is on average two times higher than K20 $T_{core}$ )."
" Above K20-100 keV/cm? the difference between K20 (T,,,,) and K20 (Telobal) Scatters around the equality.", Above K20=100 $^{2}$ the difference between K20 $T_{core}$ ) and K20 $T_{global}$ ) scatters around the equality.
" Therefore, we expect the K20-csp relation to be even steeper for the local clusters, if one uses a resolved central temperature."," Therefore, we expect the $c_{SB}$ relation to be even steeper for the local clusters, if one uses a resolved central temperature."
 We do not find any evidence for bimodality in the distributions of K20., We do not find any evidence for bimodality in the distributions of K20.
" The central cooling time is the measure most often used to quantify cool-cores, as it provides a time-frame for the evolutionary state of the gas."," The central cooling time is the measure most often used to quantify cool-cores, as it provides a time-frame for the evolutionary state of the gas."
" Adopting an isobaric cooling model for the central gas, too; can be computed as: where A(T), A(T) ng, Ne and T are the cooling function, gas number density, electron number density and temperature respectively (Peterson Fabian 2006)), and with n,=1.9n,."," Adopting an isobaric cooling model for the central gas, $t_{cool}$ can be computed as: where $\Lambda(T)$, $n_{g}$ , $n_{e}$ and T are the cooling function, gas number density, electron number density and temperature respectively (Peterson Fabian \cite{peterson05}) ), and with $n_{g}$ $n_{e}$."
" Using the global cluster temperature and thus considering our results as upper limits, we obtained the central cooling time measured within a 20 kpc radius (see Fig. 12))."," Using the global cluster temperature and thus considering our results as upper limits, we obtained the central cooling time measured within a 20 kpc radius (see Fig. \ref{tcool}) )."
" As expected, the {εουι distributions of the local and distant samples resemble the distributions of K20."," As expected, the $t_{cool}$ distributions of the local and distant samples resemble the distributions of K20."
" The local clusters span the wide range of ages [0.7 - 32.6] Gyr, whereas the RDCS+WARPS sample is limited to [4.7 - 14.3] Gyr."," The local clusters span the wide range of ages [0.7 - 32.6] Gyr, whereas the RDCS+WARPS sample is limited to [4.7 - 14.3] Gyr."
" For completeness, we report for each sample the fraction of clusters with a central cooling time lower than the age of the Universe at the cluster redshift:in the"," For completeness, we report for each sample the fraction of clusters with a central cooling time lower than the age of the Universe at the cluster redshift:in the"
"ιο)=MfarVr) AL, Vir). rs 0.1rj 0.01ri argxMy/c2. Mje, @ Mj c>7o a r2rp ry. &.. (V5) (Vi) c. daa-c) Sdd-L) M, Vf, ch MM. rove. riuiry (11)). Mp 8: >50 ~3020—30 10—15 5—8 & ÁrZ9rg. Mj.<11 c;=1-10 M,=0.5710 rg/r,=0.0571. (1)). 5 8)) a, [1]]) Mt0. ","$\rho_w(r)=\dot{M}_\ast/4\pi r^2 V_w(r)$ $\dot{M_\ast}$ $V_w(r)$ $r_s$ $0.1\,r_B$ $0.01\,r_B$ $\alpha$$r_B\propto M_p/\cs^2$ $M_p$$\cs$ $\alpha$ $M_p$ $\cs^{-2}$ $\alpha$ $r \gg r_p$ $r_p$ \ref{fig:prf2}, $V_p$ $V_w$ $\cs$ \ref{fig:prf2}a \ref{fig:prf2}d $\mach_p$ $\mach_w$ $\mach_w$ $r_{\rm over}$ $\Delta r_{\rm over}/r_p$ \ref{equ:over}) $\mach_w$ \ref{fig:prf2}: $>50$ $\sim30$$20-30$ $10-15$ $5-8$ \ref{fig:prf2} $r\gg r_B$ $\mach_w\leq11$ $\cs=1$ $\mach_p=0.5$ $r_B/r_p=0.05$ \ref{equ:aone}) \ref{fig:prf1} \ref{fig:prf2}) $\alpha_1$ \ref{equ:aone}] $\mach_w=0$ "
The largest scale of the acoustic waves was set bv the sound horizon.,The largest scale of the acoustic waves was set by the sound horizon.
 Due to the rapiditv of the decoupling process alter the universe recombined. and the lack of significant interactions therealter. these largest scale wavelengths remain in close (ο their primordial state today.," Due to the rapidity of the decoupling process after the universe recombined, and the lack of significant interactions thereafter, these largest scale wavelengths remain in close to their primordial state today."
 Such oscillations take the form of patterns on this primordial sound. horizon scale. and harmonics. in the spatial distribution of photons and barvons.," Such oscillations take the form of patterns on this primordial sound horizon scale, and harmonics, in the spatial distribution of photons and baryons."
 These acoustic waves in the photon number (or temperature) distribution were detected some 33 vears after the CMD discovery., These acoustic waves in the photon number (or temperature) distribution were detected some 33 years after the CMB discovery.
 The acoustic waves in the barvon spatial distribution were detected in 2005., The acoustic waves in the baryon spatial distribution were detected in 2005.
 The pattern in (he CMD photons shows up as peaks and troughlis of order unity deviation. making precision measurement of the angular scale of the primordial sound horizon possible with modern wide area surveys such as the WMADP satellite.," The pattern in the CMB photons shows up as peaks and troughs of order unity deviation, making precision measurement of the angular scale of the primordial sound horizon possible with modern wide area surveys such as the WMAP satellite."
 The theoretical derivation of the physical scale as a function of cosmology is straightforward due to the simple. well understood physics entering (the photon-barvon coupling and decoupling. and the linear nature of the acoustic perturbations. due to (netuation amplitudes of less than 10‘in the nunmber density. seeded (presumably) by early universe inflation.," The theoretical derivation of the physical scale as a function of cosmology is straightforward due to the simple, well understood physics entering the photon-baryon coupling and decoupling, and the linear nature of the acoustic perturbations, due to fluctuation amplitudes of less than $10^{-4}$ in the number density, seeded (presumably) by early universe inflation."
 From the two elements of the measured angular scale and theoretical physical scale. one obtains the angular distance to the decoupling epoch.," From the two elements of the measured angular scale and theoretical physical scale, one obtains the angular distance to the decoupling epoch."
 The barvon sile of the effect is similar. but with some important differences.," The baryon side of the effect is similar, but with some important differences."
 For one thine. we do not detect barvons directly in the way we detect photons.," For one thing, we do not detect baryons directly in the way we detect photons."
 Instead. we detect light emitted from processes involving barvous or electrons. or light affected in some way by the gravitational potential of mass in a structure (galaxy or cluster. sav).," Instead we detect light emitted from processes involving baryons or electrons, or light affected in some way by the gravitational potential of mass in a structure (galaxy or cluster, say)."
 While electrons should trace (he barvon pattern well. and can be neglected in the mass effects. (since a proton oulweighs an electron by some 2000 (mes). other important components of mass exist besides barvons.," While electrons should trace the baryon pattern well, and can be neglected in the mass effects (since a proton outweighs an electron by some 2000 times), other important components of mass exist besides baryons."
 Indeed. cold clark matter particles contribute six limes more (han barvons (ο the mass density of the universe., Indeed cold dark matter particles contribute six times more than baryons to the mass density of the universe.
 Thus. the spatial pattern of oscillations traced by massive structures has been diluted relative to (he primordial barvon acoustic oscillations.," Thus, the spatial pattern of oscillations traced by massive structures has been diluted relative to the primordial baryon acoustic oscillations."
 Furthermore. the barvons aller decoupling found a ready made pattern of gravitational potentials. Irom the cold dark matter. waiting to influence them.," Furthermore, the baryons after decoupling found a ready made pattern of gravitational potentials, from the cold dark matter, waiting to influence them."
 So the amplitude of the barvon acoustic oscillations is not of order one. like (he photons. but rather awzx5U..," So the amplitude of the baryon acoustic oscillations is not of order one, like the photons, but rather $\la5\%$."
 On the plus side. these oscillations are not relative to a <10| base perturbation amplitude. but rather one of order 10.1. since (he matter perturbations have been amplified by gravitational instability since the (time of decoupling.," On the plus side, these oscillations are not relative to a $<10^{-4}$ base perturbation amplitude, but rather one of order $10^{-1}$, since the matter perturbations have been amplified by gravitational instability since the time of decoupling."
 Why then isnt it üivial to detect barvon acoustic oscillations. if their absolute amplitude is so much greater than the photon acoustic peaks?," Why then isn't it trivial to detect baryon acoustic oscillations, if their absolute amplitude is so much greater than the photon acoustic peaks?"
" Unfortunately. many more CMD photons are available to be detected = 10! pass through an outstretched hand each second. while there are fewer (han 107""10 galaxies in the entire visible universe."," Unfortunately, many more CMB photons are available to be detected – $10^{15}$ pass through an outstretched hand each second, while there are fewer than $10^{10}$ galaxies in the entire visible universe."
 Furthermore. there is much," Furthermore, there is much"
strong that it is unlikely that the acceleration can proceed at the same rate wilh no change in phivsies after such a dramatic shock restructuring (pre-compression £e maxv rise by 1-2 orders of magnitudes depending on the Mach number).,strong that it is unlikely that the acceleration can proceed at the same rate with no change in physics after such a dramatic shock restructuring (pre-compression R may rise by 1-2 orders of magnitudes depending on the Mach number).
 Time dependent numerical simulations Ixangetal. 2002)) show that the modifications occur very quickly. ancl compression is increased substantially even belore 1l (note that this would be consistent with the bifurcation diagram in Figure 1. for initial OL. where is the ratio of CR number density at the shock to that of the background. plasma [ar upstream).," Time dependent numerical simulations \citealt{kjg02}) ) show that the modifications occur very quickly, and compression is increased substantially even before 1 (note that this would be consistent with the bifurcation diagram in Figure \ref{fig:bif} for initial 0.1, where is the ratio of CR number density at the shock to that of the background plasma far upstream)."
 The shock modification. in turn. must follow rather abruptly after the maximum momentum has passed through the critical value.," The shock modification, in turn, must follow rather abruptly after the maximum momentum has passed through the critical value."
 Ho was argued recently (Malkovetal.d2000)) that this should drive crucial acceleration parameters such as (he maximum monmentum injection rate back (o their critical values which must limit shock modification and settle it at some marginal level. the so called sell-organized critical (SOC) level (see also Jones2000:Malkov2001:Malkov&Diamond2001 for more discussions of the critical interrelation between (he injection. maximum energy and shock structure).," It was argued recently \citealt{mdv00}) ) that this should drive crucial acceleration parameters such as the maximum momentum and injection rate back to their critical values which must limit shock modification and settle it at some marginal level, the so called self-organized critical (SOC) level (see also \citealt{tom00, mdru01, mdFer01} for more discussions of the critical interrelation between the injection, maximum energy and shock structure)."
 \lathematically. the SOC state is characterized by (he requirement of merging of the two critical points on the bifurcation diagram in Figure l into one inflection point on the pH) eraph.," Mathematically, the SOC state is characterized by the requirement of merging of the two critical points on the bifurcation diagram in Figure \ref{fig:bif} into one inflection point on the (R) graph."
 Perhaps the most appealing aspect of this approach is its abilitv to predict the values of all three order parameters (injection rate. maxim momentum ancl compression ratio) given (he only control parameter (the Mach number) just from our knowledge of the nonlinear response Pima.) shown in Figure 1.. and no further calculations.," Perhaps the most appealing aspect of this approach is its ability to predict the values of all three order parameters (injection rate, maximum momentum and compression ratio) given the only control parameter (the Mach number) just from our knowledge of the nonlinear response ) shown in Figure \ref{fig:bif}, and no further calculations."
 IIowever. the required backreaction mechanisms on the injection and maximum momentum need to be demonstrated to operate.," However, the required backreaction mechanisms on the injection and maximum momentum need to be demonstrated to operate."
 We have already. discussed at a qualitative level how (he injection rate is reduced by shock modification., We have already discussed at a qualitative level how the injection rate is reduced by shock modification.
" The subject of this paper has been the reduction of particle momenta related to the formation of a spectral break alp,.. as a result of wave compression in a modified shock precursor."," The subject of this paper has been the reduction of particle momenta related to the formation of a spectral break at, as a result of wave compression in a modified shock precursor."
" The position of the spectral break is universally related to the degree of svstem nonlinearity /2.. since p,j,,,/A.. Ilence. the problem seems {ο be converted to the study of nonlinear properties of the acceleration (hat are formally known from the analvtie solution shown in Figure 1.."," The position of the spectral break is universally related to the degree of system nonlinearity R, since /R. Hence, the problem seems to be converted to the study of nonlinear properties of the acceleration that are formally known from the analytic solution shown in Figure \ref{fig:bif}."
 IIlowever. the injection rate that is now required for accurate determination of the spectral break through R.. may currently be obtained only [rom the SOC ansatz.," However, the injection rate that is now required for accurate determination of the spectral break through R, may currently be obtained only from the SOC ansatz."
" It should be also mentioned that strong reduction of is obviously not to be expected in oblique shocks. where the resonance relation D is approximately preserved due to the compression of D simultaneously with hr. An equally important problem is that strong losses of particles between and must slow down the growth of p4,0) due to the reductionof resonant waves."," It should be also mentioned that strong reduction of is obviously not to be expected in oblique shocks, where the resonance relation B is approximately preserved due to the compression of B simultaneously with k. An equally important problem is that strong losses of particles between and must slow down the growth of (t) due to the reductionof resonant waves."
 As we argued in section 2.1.. this may result in," As we argued in section \ref{sec:signif}, , this may result in"
a function of 24um flux is not an artificial selection effect due to the relative depths of our MIPS and r*-band COSMOS observations.,a function of $\mu$ m flux is not an artificial selection effect due to the relative depths of our MIPS and $r^+$ -band COSMOS observations.
 Even at the 24um flux limit of mmJy the optical COSMOS data are deep enough to identify galaxies as red as the color threshold used in the criterion proposed by Deyetal. (2008)., Even at the $\mu$ m flux limit of mJy the optical COSMOS data are deep enough to identify galaxies as red as the color threshold used in the criterion proposed by \citet{Dey:08}.
. Also it implies that the of the initial 24um detections that we could not identify at optical and near-IR wavelengths (Sect.2) satisfy Eq., Also it implies that the of the initial $\mu$ m detections that we could not identify at optical and near-IR wavelengths (Sect.2) satisfy Eq.
 4 independently of their mid-IR flux density., \ref{eq:DOG} independently of their mid-IR flux density.
 Since these extremely red colors are unlikely associated to galaxies at z«1 (see reffig:histop OG1aa)therelativecontributiono f galaxiesselectedat1. DiZi ;rioncouldreach~45% at P54um0.08 mmjJy., Since these extremely red colors are unlikely associated to galaxies at $z<1$ (see \\ref{fig:histo_DOG1}a a) the relative contribution of galaxies selected at $1.5<z<3$ with this criterion could reach $\sim$ at $F_{24\mu m} > 0.08$ mJy.
 As we already argued this selection of high-redshift galaxies based on extremely red mid-IR/optical colors is by construction biased toward dust-obscured sources with very faint optical luminosities., As we already argued this selection of high-redshift galaxies based on extremely red mid-IR/optical colors is by construction biased toward dust-obscured sources with very faint optical luminosities.
 For example we found that of these objects have rt> 26mmag while the r*—band magnitude distribution of the MIPS sources with 1.5«z2.8 peaks at ΤΙ~ 25mmag (only of them have rt> 26mmag)., For example we found that of these objects have $r^+>26$ mag while the $r^+$ –band magnitude distribution of the MIPS sources with $1.5 < z < 2.8$ peaks at $r^+ \sim 25$ mag (only of them have $r^+>26$ mag).
" Although they do not represent a dominant population in terms of source density, this nicely illustrates the contribution of luminous high-redshift star-forming galaxies that may be missed by optical surveys because of dust obscuration."," Although they do not represent a dominant population in terms of source density, this nicely illustrates the contribution of luminous high-redshift star-forming galaxies that may be missed by optical surveys because of dust obscuration."
" In fact, of the Optically-Faint IR-bright sources have an extinction of E(B—V)>0.4 according to the optical SED fitting of Ilbertetal.(2009),, while this percentage decreases to only for the whole population of MIPS sources in the same redshift range."," In fact, of the Optically-Faint IR-bright sources have an extinction of $E(B-V) \ge 0.4$ according to the optical SED fitting of \citet{Ilbert:09}, while this percentage decreases to only for the whole population of MIPS sources in the same redshift range."
" Furthermore, galaxies selected with this technique differ dramatically from those identified with the BM/BX selection, which we found to be strongly biased against dusty high-redshift galaxies."," Furthermore, galaxies selected with this technique differ dramatically from those identified with the BM/BX selection, which we found to be strongly biased $against$ dusty high-redshift galaxies."
 In BzKreffig: giagramwecompareintheBzKdiagramthedistributto Faint R—brightsourceswiththedistributionof estinodi rheior, In \\ref{fig:BzK_diagram} we compare in the $BzK$ diagram the distribution of the Optically-Faint IR-bright sources with the distribution of galaxies selected with our modified BM/BX criterion.
" overlapbe samplesisrelativelysmallasonly8%oftheM IP Ssourcesatl. 2.8satis fythetwoselectioncriteria, FaintI R—brightobj ectsarecharacterizedbymuchreddercolor, "," The overlap between the two sub-samples is relatively small as only of the MIPS sources at $1.5 < z < 2.8$ satisfy the two selection criteria, and not surprisingly the Optically-Faint IR-bright objects are characterized by much redder colors than the BM/BX sources."
"ziip brightgalariesandtheDM/B.X μα tec 2in the BzK diagram selectedsourcesshareverysintfartoverlapbetweenthesetwopopulaftreffig:BzK giagram)), thelackof lumingsi οiverhi"," Since the Optically Faint IR-bright galaxies and the BM/BX selected sources share very similar redshift distributions (compare 2a 3b) and given the extinction vector at $z \sim 2$ in the $BzK$ diagram \\ref{fig:BzK_diagram}) ), the lack of overlap between these two populations is mostly due to the effect of dust extinction."
 In reffig:overlap we illustrate in a more quantitative way how the Optically Faint IR-bright objects and the other selections overlap with each other., In \\ref{fig:overlap} we illustrate in a more quantitative way how the Optically Faint IR-bright objects and the other selections overlap with each other.
" Since we found that the high efficiency of the BzK criterion is maintained up to z~3 (at least when applied to our mid-IR sample, see 44.1), we compared our different selections over the largest possible redshift range (ie., 1.5«z 2.8) to minimize the statistical uncertainties."," Since we found that the high efficiency of the $BzK$ criterion is maintained up to $z \sim 3$ (at least when applied to our mid-IR sample, see 4.1), we compared our different selections over the largest possible redshift range (i.e., $1.5 < z < 2.8$ ) to minimize the statistical uncertainties."
 More than half of the population of Optically Faint IR-bright sources are also selected as IRAC Peakers indicating that an important part of the sources presenting extreme MIR color excess have an SED dominated by the stellar bump., More than half of the population of Optically Faint IR-bright sources are also selected as IRAC Peakers indicating that an important part of the sources presenting extreme MIR color excess have an SED dominated by the stellar bump.
" In the two previous sections we described a number of color selection techniques that have been widely used for identifying star-forming galaxies at 1.5«z3, and pu respective NAM to the total galawiesselecte t theirου| απλο sources ineenthetwo. withourdust oe desections’ suffer andno"," In the two previous sections we described a number of color selection techniques that have been widely used for identifying star-forming galaxies at $1.5<z<3$, and we quantified their respective contribution to the total number density of MIPS-selected high-redshift sources in order to estimate the bias that these selections suffer because of dust extinction."
"tsurprisinglyth& pico,co extinction. We want to extend this sth th μμ) ]ή we onstcefmiMS sd dsht Hn IR Hannonuseit ford reipnd.the cosmic Sum detai"," We want to extend this analysis by measuring the mid-IR luminosity function of the different sub-samples of galaxies selected based on these techniques, so as to infer their contribution to the total IR luminosity density of the Universe and the cosmic star formation rate density observed at $z \sim 2$."
"l the methods used to obtain rest-frame Luminosity Function (LF), the IR luminosity density and the Star Formation Rate (SFR)."," Here we detail the methods used to obtain the $\mu$ m rest-frame Luminosity Function (LF), the IR luminosity density and the Star Formation Rate (SFR)."
 The galaxy luminosity functions were all computed at rest-frame 8jm. The redshift distributions of the sources selected with the techniques described earlier mostly peak at z~2 (Figs.22a 3b) and at these redshifts the measure of 84m luminosities (Lgm) based on 24m fluxes is therefore almost independent of the SED templates assumed for deriving the k-corrections., The galaxy luminosity functions were all computed at rest-frame $\mu$ m. The redshift distributions of the sources selected with the techniques described earlier mostly peak at $z \sim 2$ 2a 3b) and at these redshifts the measure of $\mu$ m luminosities $L_{8\mu m}$ ) based on $\mu$ m fluxes is therefore almost independent of the SED templates assumed for deriving the $k$ -corrections.
" Also, the 84m"," Also, the $\mu $ m"
The first 5 are observed SEDs aud SEDs 6 and 7 are calculated from the models of (2003).,The first 5 are observed SEDs and SEDs 6 and 7 are calculated from the models of \citet{bru03}.
. The properties of these models are given in Table 1.., The properties of these models are given in Table \ref{tab-temp}.
 Template 6 is for a 50 million vear old galaxy and Template 7 is for a 10 million vear old galaxy. with the lowest metallidtv available in (he Bruzual&Charlot(2003) models., Template 6 is for a 50 million year old galaxy and Template 7 is for a 10 million year old galaxy with the lowest metallicity available in the \citet{bru03} models.
 Template 7 is similar (0 a Pop., Template 7 is similar to a Pop.
 HI SED but redder than a true Pop., III SED but redder than a true Pop.
 HI SED (see Schaerer(2002) for the expected differences)., III SED (see \citet{sch02} for the expected differences).
 Although a Pop., Although a Pop.
 HI SED may be bluer than template 7. the sharp increase in the 3.6; to 1.65an fluctuation power at high redshift is due to the Lyman break passing through the 1.6j/mi and is therefore relatively independent of the SED to the red of the break.," III SED may be bluer than template 7, the sharp increase in the $\micron$ to $\micron$ fluctuation power at high redshift is due to the Lyman break passing through the $\micron$ and is therefore relatively independent of the SED to the red of the break."
 Thompsonetal.(2006) interpolated between the SEDs in increments of 0.1 to produce a final set of 61 different model SEDs., \citet{thm06} interpolated between the SEDs in increments of 0.1 to produce a final set of 61 different model SEDs.
 Figure 2. shows the distribution of the template numbers for all sources in the NUDF with redshifts greater than 2.0., Figure \ref{fig-temp} shows the distribution of the template numbers for all sources in the NUDF with redshifts greater than 2.0.
 The predominance of sources concentrate on templates 6 aud 7. therefore. we will take them as the most probable SEDs for sources fainter than our source detection limit.," The predominance of sources concentrate on templates 6 and 7, therefore, we will take them as the most probable SEDs for sources fainter than our source detection limit."
 We use these SEDs and the effective response functions in the NICMOS and IRAC filters to caleulate the expected ratios of the power at redshifts between 0 and 15., We use these SEDs and the effective response functions in the NICMOS and IRAC filters to calculate the expected ratios of the power at redshifts between 0 and 15.
 Absorption in the Inter-Galactie Medium is modeled according to the treatinent eiven in Macauetal.(1996)., Absorption in the Inter-Galactic Medium is modeled according to the treatment given in \citet{mad96}.
. The horizontal lines in Figure 2. show (he observed ratios of the fluctuation power in the all source subtracted images using: (Bis analvsis. INashlinskyv.etal.(2007). ancl etal. (2007).," The horizontal lines in Figure \ref{fig-col} show the observed ratios of the fluctuation power in the all source subtracted images using; this analysis, \citet{kas07b} and \citet{thm07}."
". Since the excess fluctuation power attributed to high redshift galaxies by ]xashlinskvetal.(2007b) is at angular scales greater than 30"" the horizontal lines show the ralios observed αἱ 100"". the largest NICMOS angular scale."," Since the excess fluctuation power attributed to high redshift galaxies by \citet{kas07b} is at angular scales greater than $30\arcsec$ the horizontal lines show the ratios observed at $100\arcsec$, the largest NICMOS angular scale."
 The observed {μι to L.6;an fluctuation power ratio is consistent with the lower redshift template G predictions. but is incompatible with the caleulated ratios al 2>8 for either of the templates.," The observed $\micron$ to $\micron$ fluctuation power ratio is consistent with the lower redshift template 6 predictions, but is incompatible with the calculated ratios at $z>8$ for either of the templates."
 This is strong evidence that the source subiracted fIuctuations in the NICAIOS images are due to galaxies al 2«8., This is strong evidence that the source subtracted fluctuations in the NICMOS images are due to galaxies at $z<8$.
 The observed 3.65 to 1.670n ratio lies between the ratios predicted bv templates 6 aud 7., The observed $\micron$ to $\micron$ ratio lies between the ratios predicted by templates 6 and 7.
 Given the good match to the NICMOS flux ratios for template 6 we would expect higher 3.6/an. fluctuations than observed., Given the good match to the NICMOS flux ratios for template 6 we would expect higher $\micron$ fluctuations than observed.
 The difference may be due to cosmic variance since the fields are different., The difference may be due to cosmic variance since the fields are different.
 It may also be that the much broader PSFs subiract more background [αν than the IST PSEs., It may also be that the much broader PSFs subtract more background flux than the HST PSFs.
 For 2>]2 the predicted ratios become incompatible with observations., For $z>12$ the predicted ratios become incompatible with observations.
 The 4.5jan to 3.61 ratio is essentially independent of redshift at the redshilts considered ancl therefore does nol contribute any information on the redshift epoch of the Hactuations., The $\micron$ to $\micron$ ratio is essentially independent of redshift at the redshifts considered and therefore does not contribute any information on the redshift epoch of the fluctuations.
 The apparent agreement between the predicted ancl observed ratios Dor those bands. however. is further evidence that the observed. fInctuations are due to faint galaxies at redshifts less than 3.," The apparent agreement between the predicted and observed ratios for those bands, however, is further evidence that the observed fluctuations are due to faint galaxies at redshifts less than 8."
The global photometric properties of normal galaxies can be summarized. by the folowing parameters: (1) absolute magnitude. (2) colours. (3) compactess. and (4) raclia distribution and scale-Iengh.,"The global photometric properties of normal galaxies can be summarized by the following parameters: (1) absolute magnitude, (2) colours, (3) compactness, and (4) radial distribution and scale-length."
 To a firs apooximation. high-luminosity giant galaxies are either early-type (Illiptica or SO) and have a radial profile that is well-lit by a. de Vaucouleurs (1948) ΕΣ Jaw or late-ty»e (spiral or irregular) and have a racial. profile hat is well-fit » an exponentia law (Freeman 1970).," To a first approximation, high-luminosity giant galaxies are either early-type (Elliptical or SO) and have a radial profile that is well-fit by a de Vaucouleurs (1948) $r^{1/4}$ law or late-type (spiral or irregular) and have a radial profile that is well-fit by an exponential law (Freeman 1970)."
 Lower luminosity galaxies (dwarl irregulars and dwarf spheroidals) also have exponential ligh profiles (Bingeeli Cameron 1991) but have lower surface brightnesses., Lower luminosity galaxies (dwarf irregulars and dwarf spheroidals) also have exponential light profiles (Binggeli Cameron 1991) but have lower surface brightnesses.
 Many late-twpe galaxies have light profiles which can be decomposed into a bulge Get) and. disk (exponential) part. (IXormendy. 1977. Went 1985).," Many late-type galaxies have light profiles which can be decomposed into a bulge $r^{1/4}$ ) and disk (exponential) part (Kormendy 1977, Kent 1985)."
 More subtle cllects can be seen by. examining correlations between the various parameters for a sample of galaxies (see Fig., More subtle effects can be seen by examining correlations between the various parameters for a sample of galaxies (see Fig.
 1): Ipt (1) Eary type galaxies are τοσο ancl have older stars than [ate-ty2ο OLlCS., 1): 1pt (1) Early type galaxies are redder and have older stars than late-type ones.
 1pt (2) Luminous early-type galaxies tend to have higher surlace-brigitnesses than luminous late-tvpe galaxies., 1pt (2) Luminous early-type galaxies tend to have higher surface-brightnesses than luminous late-type galaxies.
 They are more compact (there do exist. however. late-tvpe galaxies with compact cores): Ipt (3) Late-twvpe galaxies tend to have lower surface-brghtnesses as their luminosity decreases.," They are more compact (there do exist, however, late-type galaxies with compact cores); 1pt (3) Late-type galaxies tend to have lower surface-brightnesses as their luminosity decreases."
 This trene continues into the regime of cdwarl galaxies: Ipt (4) Vhere is a tendency for more luminous carly- galaxies to have lower surlacc-brightnesses., This trend continues into the regime of dwarf galaxies; 1pt (4) There is a tendency for more luminous early-type galaxies to have lower surface-brightnesses.
 This. is characterized by the familiar fundamental plane for elliptical galaxies (Ixormendy Djorgovski 1989): Ipt, This is characterized by the familiar fundamental plane for elliptical galaxies (Kormendy Djorgovski 1989); 1pt
Iuerstellar shocks generated by violent stellar activities. such as supernova explosions aud »otostelar outflows. have profouud ellects ou the surrounding interstellar medium.,"Interstellar shocks generated by violent stellar activities, such as supernova explosions and protostellar outflows, have profound effects on the surrounding interstellar medium."
 Shocks propagating into dense molecular clouds can heat the gas to several hiuxlred or several thousand Kelvin aud »odi ‘ich spectra in the infrared. spectral regiou., Shocks propagating into dense molecular clouds can heat the gas to several hundred or several thousand Kelvin and produce rich spectra in the infrared spectral region.
 Fast shocks| with speeds larger than £0 και 1 a σαν dissociative: they destroy. molecules aud lonize atoms. generating strong atomic lne-structure emissious.," Fast shocks with speeds larger than 40 km $^{-1}$ are usually dissociative; they destroy molecules and ionize atoms, generating strong atomic fine-structure emissions."
 Ou the other haud. most molecules swvive in slow shocks.," On the other hand, most molecules survive in slow shocks."
 Heating excites various iiolecular species via collisional processes. causiug he gas to glow.," Heating excites various molecular species via collisional processes, causing the gas to glow."
 A large number of nolecula ‘line features lave been observed iu. associatiou with shock-allected areas. includiug cooling lines fro1 He. CO. HD aud HeO. Due to its abiliy to reveal species that are difficult to observe within cold quiesceu eas. a shock wave serves as a “searchlieht™ for the pliysical structure of molecular cloucls.," A large number of molecular line features have been observed in association with shock-affected areas, including cooling lines from ${_2}$, CO, HD and ${_2}$ O. Due to its ability to reveal species that are difficult to observe within cold quiescent gas, a shock wave serves as a “searchlight” for the physical structure of molecular clouds."
 Moreover. shocks alter tie Chemical composition of gas by criving mauy eucothermic reactious. oue of which is le Conversion of para molecular hydrogen to ortho hydrogeu.," Moreover, shocks alter the chemical composition of gas by driving many endothermic reactions, one of which is the conversion of para molecular hydrogen to ortho hydrogen."
 La previous studies. it has been found that marv sources exhibit Πο ortho-to-para ratio (hereafter OPR) markedly ess than the equilibrium value ~3 (Neufeld et al.," In previous studies, it has been found that many sources exhibit ${_2}$ ortho-to-para ratio (hereafter OPR) markedly less than the equilibrium value $\sim 3$ (Neufeld et al."
 1998: Cabrit. et al.," 1998; Cabrit, et al."
 1999: Neufeld et al., 1999; Neufeld et al.
 2006. rerealter 06: Nettled οἱ al.," 2006, hereafter N06; Neufeld et al."
 2007. hereafter NOT).," 2007, hereafter N07)."
 Adoptiug a two component tjoceel containiug a mixture of wari aid hot gas. NOG aud NOT found that [or the sources we are stuclying in this yaper — ICLI8C. W28. WEI. 8039)|. HH? and HH51- the OPR values in the wart1 gas colnponelit CP~300—600 IX) a'e 2.12. 0.93. 1.58. 0.65. Q.21 and 0.11-0.18 respectively.," Adopting a two component model containing a mixture of warm and hot gas, N06 and N07 found that for the sources we are studying in this paper – IC443C, W28, W44, 3C391, HH7 and HH54 – the OPR values in the warm gas component $T \sim 300-600$ K) are 2.42, 0.93, 1.58, 0.65, 0.21 and 0.41–0.48 respectively."
 They proposed tliat he non-LTE OPR valtes obtainec may correspond to the temperature at ai earier epoch. owing o the low efficiency of j)ara-to-ortho conversion in uou-dissociative shocks: this then can provide is useful information o1 the evolujon titnescale.," They proposed that the non-LTE OPR values obtained may correspond to the temperature at an earlier epoch, owing to the low efficiency of para-to-ortho conversion in non-dissociative shocks; this then can provide us useful information on the evolution timescale."
 In this paper we luvestigate umbolecular shocks associated with four supernova remuants — ICIISC. ος. WELE ane 3C391 aix two Herbig-Haro objects - HH7 aud HH51.," In this paper we investigate molecular shocks associated with four supernova remnants – IC443C, W28, W44 and 3C391 and two Herbig-Haro objects – HH7 and HH54."
 All six sources have extensive evidence for iieractiou witl uiolecular clouds provided by imulti-waveleugth observatious., All six sources have extensive evidence for interaction with molecular clouds provided by multi-wavelength observations.
 A brief description of t1056 SOLICes Is eiven below., A brief description of these sources is given below.
" The four bright sperhova remiut sources [οο. W28. WIL aud 3€C391 have provided excellent laboratories fo ""the sttον oft heinteraction between SNR shocks and surrouxcing molecular eas."," The four bright supernova remnant sources IC443, W28, W44 and 3C391 have provided excellent laboratories for the study of the interaction between SNR shocks and surrounding molecular gas."
" W28. WEE and 3C391 are protetypes ¢ Mule ""mixed-morphiology class. whose centrally coucentrated X-ray morphologies contrast with the sheI-like radio emission (Rho et al."," W28, W44 and 3C391 are prototypes of the “mixed-morphology” class, whose centrally concentrated X-ray morphologies contrast with the shell-like radio emission (Rho et al."
 1991: RIo Peter 1996: Rho Peter 1998)., 1994; Rho Peter 1996; Rho Peter 1998).
 The radio aud X-ΑΝ NOPshology of ICE13 also shows similarities lo the mixed- class. although wibh aciditioral X-ray components.," The radio and X-ray morphology of IC443 also shows similarities to the mixed-morphology class, although with additional X-ray components."
 X-ray. observatious show the four 'eranants to be filled with a large amount of hot gas. with deusity »~ L1 — 10 E” accordiug o tlie radiative model of Harrus et al. (," X-ray observations show the four remnants to be filled with a large amount of hot gas, with density $n~\sim$ 1 – 10 $^{-3}$ according to the radiative model of Harrus et al. ("
1997) and. Chevalier (1999): this iuplicates au interaction with relatively deuse environments — probably intercloucl gas.,1997) and Chevalier (1999); this implicates an interaction with relatively dense environments | probably intercloud gas.
 More convincing evidence for the interaction with deuse clouds arises directly. from the detection of various excited. molecules at onger wavelengths., More convincing evidence for the interaction with dense clouds arises directly from the detection of various excited molecules at longer wavelengths.
 As stummarized by Reach et al. (, As summarized by Reach et al. (
2005) aud NOT for eacl individual source,2005) and N07 for each individual source
In order to estimate (he brightness of the reported companion. a model for à 50.7; brown cwarf at 5 Gyr was chosen.,"In order to estimate the brightness of the reported companion, a model for a $50M_{\rm J}$ brown dwarf at 5 Gyr was chosen."
 This mass lies conservatively in the lower range of possible values and required no interpolation within available models., This mass lies conservatively in the lower range of possible values and required no interpolation within available models.
 The model age is likely to be greater than the 3.67 Gyr cooling age (Bergeron.Leggett.&Ruiz2001) of van Maanen 2 plus the ~0.5 Gyr main sequence lifetime (Maeder1989). lor a ~4A. progenitor of the 0.82... white cdwarl (Weidemann1937.1990.2000:Dragaglin.Renzinial. 2001).," The model age is likely to be greater than the 3.67 Gyr cooling age \citep*{ber01}
 of van Maanen 2 plus the $\sim0.5$ Gyr main sequence lifetime \citep
{mae89} for a $\sim4M_{\odot}$ progenitor of the $0.83M_{\odot}$ white dwarf \citep{wei87,wei90,wei00,bra95,ber01}."
.. A substellar companion of this mass ancl age would be a late T dwarf (Zong~800 IX) and have Mj;=143 mag which is L—12.5 mag at 4.4 pe (Burrowsetal.1997)., A substellar companion of this mass and age would be a late T dwarf $T_{\rm eff}\sim800$ K) and have $M_{L^{'}}=14.3$ mag which is $L^{'}=12.5$ mag at 4.4 pc \citep{bur97}.
. Photometric L' band data on known brown dwarls do exist and the measurements agree with (he models used here to within 0.3 mag for spectral tvpe T6 (Leeeettetal.2002).., Photometric $L^{'}$ band data on known brown dwarfs do exist and the measurements agree with the models used here to within 0.3 mag for spectral type T6 \citep{leg02}. .
 Van Maanen 2 is predicted to have L'=11.43 mag based on the model predicted V.—A=0.84 color of a 6750 Ix helium white dwarl (Bergeron.Wesemael.&Beauchamp1995) and the N-—L--042 color of a 6750 IX blackbody., Van Maanen 2 is predicted to have $L^{'}=11.43$ mag based on the model predicted $V-K=0.84$ color of a 6750 K helium white dwarf \citep{ber95} and the $K-L^{'}=0.12$ color of a 6750 K blackbody.
 V.—ΑΝ=0.89 is the measured color — the extrapolation was done [rom V. in case of any contamination by a companion at A., $V-K=0.89$ is the measured color – the extrapolation was done from $V$ in case of any contamination by a companion at $K$.
 The reduced L’ image (Figure 1)) shows no indication whatsoever of a companion with the brightness expected from a brown dwarf of the twpe reported bv Makarov(2004)., The reduced $L^{'}$ image (Figure \ref{fig1}) ) shows no indication whatsoever of a companion with the brightness expected from a brown dwarf of the type reported by \citet{mak04}.
".. Frou the published orbital parameters. the companion should have been at à separation of 0.19"" and position angle of 274° on the date of the observation."," From the published orbital parameters, the companion should have been at a separation of $0.19''$ and position angle of $274\arcdeg$ on the date of the observation."
" The full width at half maximum ol van Maanen 2 in the reduced image is 0.089"" and the distance to the first Airy ring is z0.14"".", The full width at half maximum of van Maanen 2 in the reduced image is $0.089''$ and the distance to the first Airy ring is $\approx0.14''$.
 There are two extremely faint features within the Airy ring at position angles of 284 and 2957., There are two extremely faint features within the Airy ring at position angles of 284 and $295\arcdeg$.
 Small aperture flux measurements. relative to the primary. al eielt clilferent evenly spaced locations around the Airy ring indicate Chat these features are unlikely to be real.," Small aperture flux measurements, relative to the primary, at eight different evenly spaced locations around the Airy ring indicate that these features are unlikely to be real."
 The flux at 284 and 2957. f/fy=0.069 and 0.073 respectively. are both within 2c (0.018) of the average [lux (f/fp= 0.056) in the ring and are almost certainly. artifacts due to imperfect optical guide star corrections.," The flux at 284 and $295\arcdeg$, $f/{f_0}=0.069$ and 0.073 respectively, are both within $2\sigma$ (0.018) of the average flux $f/{f_0}=
0.056$ ) in the ring and are almost certainly artifacts due to imperfect optical guide star corrections."
 If the feature al 295° were real. ils brightness alter subtracting the flux of the primary in the Airy ring implies L'=15.9 mag and a mass ol 15A. assuming an age of 5 Gyr (Duirowsοἱal.1997).," If the feature at $295\arcdeg$ were real, its brightness after subtracting the flux of the primary in the Airy ring implies $L^{'}=15.9$ mag and a mass of $\sim15M_{\rm J}$, assuming an age of 5 Gyr \citep{bur97}."
.. This is simply not massive enough to cause the reported astrometric wobble., This is simply not massive enough to cause the reported astrometric wobble.
" An artilicial star was planted at 0.19"" from the primary in order to test the ability to detect faint companions at this separation.", An artificial star was planted at $0.19''$ from the primary in order to test the ability to detect faint companions at this separation.
 The adaptive optics PSF. extracted. [rom the reduced image of van Maanen 2. was used lor the stellar profile of the planted star.," The adaptive optics PSF, extracted from the reduced image of van Maanen 2, was used for the stellar profile of the planted star."
 To be conservative. this artificial staris a [ull 2.0 magnitudes fainter than the white dwarl — this was confirmed. by placing the star al many posilions on (he image and measuring its flux," To be conservative, this artificial staris a full 2.0 magnitudes fainter than the white dwarf – this was confirmed by placing the star at many positions on the image and measuring its flux"
yaper. we assume that ry in eq. (9)),"paper, we assume that $\eta_0$ in eq. \ref{sigma}) )"
" does not vary with gas surface density Al, (see section 3).", does not vary with gas surface density $M_g$ (see section 3).
 The actor Jp/(Ld-yo) is what determines the critical column deusity where the HJ—HII transition occurs Le. the columu density where the steep slope of tlie second region starts.," The factor ${J_L/(1+\eta_0)}$ is what determines the critical column density where the ${HI\rightarrow
HII}$ transition occurs i.e. the column density where the steep slope of the second region starts."
 The rapidity of he decrease of Nyy; in the second region depeuds also on metal abundances., The rapidity of the decrease of $N_{HI\perp}$ in the second region depends also on metal abundances.
 When Z increases above 0.05 Z. metal line cooling becomes quite importaut in that region aud the transitiou gets ess sharp as Z increases., When $Z$ increases above 0.05 $Z_{\odot}$ metal line cooling becomes quite important in that region and the transition gets less sharp as $Z$ increases.
 The shape of these curves is quite important because it determines the slope of the distribution fuuctioa around the Lyman-Liuit region., The shape of these curves is quite important because it determines the slope of the distribution function around the Lyman-Limit region.
 Several papers on the HI columu deusity distribution of absorbers (Tytler1987: have shown that a power law Nap fits the data from LOM 7 to 107} em7 approximately. but not well enough to satisly statistical tests such as the test.," Several papers on the HI column density distribution of absorbers \citep{tyt87,lwt91,pet93} have shown that a power law $N_{HI}^{-1.5}$ fits the data from $10^{13}$ $^{-2}$ to $10^{21}$ $^{-2}$ approximately, but not well enough to satisfy statistical tests such as the Kolmogorov-Smirnov test."
" La this paper we shall use the deviations of Nyy; distribution (rom a power law to determine .X. after assuming that the total gas columu ceusity Nyy has a power law cistributiou [Function of the form Iu order to derive the μι distribution from g(Ny,) we must orient the absorbers raudoinmly iu the plane of the sky. assume an axial ratio //HR lor the slab. and apply the αμ—μι conversion factor."," In this paper we shall use the deviations of $N_{HI}$ distribution from a power law to determine $X$, after assuming that the total gas column density $N_H$ has a power law distribution function of the form In order to derive the $N_{HI}$ distribution from $g(N_{H\perp})$ we must orient the absorbers randomly in the plane of the sky, assume an axial ratio $h/R$ for the slab, and apply the $N_{H\perp}-N_{HI\perp}$ conversion factor."
 For our inodel fitting we shall use Z=0.02Z.. .X iuxlepencdent of redshliift aud au axial ratio ol h/It= 0.2. since clouds are likely to be neither spherical nor thin disks.," For our model fitting we shall use $Z=0.02
Z_\odot$, $X$ independent of redshift and an axial ratio of $h/R=0.2$ , since clouds are likely to be neither spherical nor thin disks."
 Results essentially do not depend on h/HR for h/HRx0.5., Results essentially do not depend on $h/R$ for $h/R\le 0.5$.
 We derive a and X by comparing the resulting /(V5) lor +=1.0 with the data present in our compilation at Nyy>1.6x104 7., We derive $\alpha$ and $X$ by comparing the resulting $f(N_{HI})$ for $\gamma=1.0$ with the data present in our compilation at $N_{HI}\ge 1.6\times 10^{17}$ $^{-2}$.
 AN is set by the normalization condition for g. based ou the observed number of absorbers.," $K$ is set by the normalization condition for $g$, based on the observed number of absorbers."
 We emphasize that it is uot possible to present the data relative to the distribution functiou iu a model independent way. due to LLS with =undetermiued τε.," We emphasize that it is not possible to present the data relative to the distribution function in a model independent way, due to LLS with undetermined $\tau_{LL}$."
 A deterioration of the available data wight result from the operation of bimuine in Ny; in order to reucder straight[οανα the comparisou with the moclel distributions., A deterioration of the available data might result from the operation of binning in $N_{HI}$ in order to render straightforward the comparison with the model distributions.
 Iustead of binning the data. we match the model cistribution to the individual detectious.," Instead of binning the data, we match the model distribution to the individual detections."
 The details of the fitting procedure are given in Paper IE: we uuderline here he main characteristics., The details of the fitting procedure are given in Paper II; we underline here the main characteristics.
 The procedure is rather similar to that used by Storrie-Lombardi.MeMahon (1996).. i iinplementing the algoritlun in order to take into consideration also the uncertainty in the determination of any single value of μις," The procedure is rather similar to that used by \citet{sto96}, but implementing the algorithm in order to take into consideration also the uncertainty in the determination of any single value of $N_{HI}$."
 Large observational errors are includect yy leaving undeterimined the “real” position of each event., Large observational errors are included by leaving undetermined the “real” position of each event.
 We determine a aud Ὃν a Maxintun Likelihood analysis to the projected HI column ceusity distribution f lixiug the real position ofeach event to the measured value of Αι when this is available: otherwise we use in the Likelihoocd, We determine $\alpha$ and $X$ by a Maximum Likelihood analysis to the projected HI column density distribution $f$ fixing the real position ofeach event to the measured value of $N_{HI}$ when this is available; otherwise we use in the Likelihood
as a function of K.R.,as a function of $k_z R$.
" As expected, the maximum growth rate decreases as KR increases, until absolute stability of the thermal continuum is achieved for a critical longitudinal wavenumber."," As expected, the maximum growth rate decreases as $k_z R$ increases, until absolute stability of the thermal continuum is achieved for a critical longitudinal wavenumber."
 This critical wavenumber is in perfect agreement with Equation (31))., This critical wavenumber is in perfect agreement with Equation \ref{eq:kzcrit}) ).
 The result for the three different temperature profiles does not show significant differences., The result for the three different temperature profiles does not show significant differences.
" We must mention that the various parametrizations of the radiative regime Prominence-1, which aim to represent different optical thicknesses of the cool prominence plasma, are not relevant for the instability of the continuum."," We must mention that the various parametrizations of the radiative regime Prominence-1, which aim to represent different optical thicknesses of the cool prominence plasma, are not relevant for the instability of the continuum."
" The thermal spectrum in the cool part of the equilibrium. namely zone 1, is always stable independently of the considered parametrization (indicated by different line styles in Fig. 2))"," The thermal spectrum in the cool part of the equilibrium, namely zone 1, is always stable independently of the considered parametrization (indicated by different line styles in Fig. \ref{fig:continuum}) )"
 In Figure 4. we have plotted the thermal continuum growth rate for k.R=1077 using the Klimehuk-Raymond radiative loss function., In Figure \ref{fig:continuumklim} we have plotted the thermal continuum growth rate for $k_z R = 10^{-2}$ using the Klimchuk-Raymond radiative loss function.
" By comparing Figures 2((a) and 4., we see that slightly larger values of the growth rate are obtained using Klimchuk-Raymond's fit with respect to the values for Hildner's fit."," By comparing Figures \ref{fig:continuum}( (a) and \ref{fig:continuumklim}, we see that slightly larger values of the growth rate are obtained using Klimchuk-Raymond's fit with respect to the values for Hildner's fit."
" However, the qualitative behavior of the continuum is similar is both cases."," However, the qualitative behavior of the continuum is similar is both cases."
" Here, we study how the unstable part of the thermal continuum is modified when perpendicular thermal conduction and magnetic diffusion are taken into account."," Here, we study how the unstable part of the thermal continuum is modified when perpendicular thermal conduction and magnetic diffusion are taken into account."
" Unless otherwise stated, Hildner's parametrization for the radiative loss function is used in all the following computations."," Unless otherwise stated, Hildner's parametrization for the radiative loss function is used in all the following computations."
" First, we set 7=0 and focus our investigation on the effect of perpendicular thermal conduction."," First, we set $\eta=0$ and focus our investigation on the effect of perpendicular thermal conduction."
 We would like to stress that the real physical value of κι given by Equation (7)) is used in the following computations., We would like to stress that the real physical value of $\kappa_\perp$ given by Equation \ref{eq:ionskappa}) ) is used in the following computations.
 As stated by VanderLinden&Goossens(1991).. the thermal continuum is replaced by a set of discrete modes when κιz0.," As stated by \citet{vanderlinden91}, the thermal continuum is replaced by a set of discrete modes when $\kappa_\perp \ne 0$."
" The eigenfunctions of these solutions display laree variations in a region surrounding the position of the thermal continuum singularity for aK,=0.", The eigenfunctions of these solutions display large variations in a region surrounding the position of the thermal continuum singularity for $\kappa_\perp = 0$.
" Figure 5. displays the temperature perturbation, 71. of the four most unstable modes of our equilibrium in the case of the sinusoidal temperature profile with //R=1. K.R=107! and m=0."," Figure \ref{fig:eigen1} displays the temperature perturbation, $T_1$, of the four most unstable modes of our equilibrium in the case of the sinusoidal temperature profile with $l/R = 1$, $k_z R = 10^{-1}$ and $m=0$."
" For simplicity, Figure 5. shows the temperature perturbation only, because Τι is between one and two orders of magnitude larger than the other perturbations."," For simplicity, Figure \ref{fig:eigen1} shows the temperature perturbation only, because $T_1$ is between one and two orders of magnitude larger than the other perturbations."
" This means that the temperature perturbation is the dominant disturbance related to the thermal modes, although these solutions produce also velocity and magnetic field perturbations."," This means that the temperature perturbation is the dominant disturbance related to the thermal modes, although these solutions produce also velocity and magnetic field perturbations."
" Figure 5. focuses on the region where the eigenfunctions show significant variations, 1.c., the conductive layer described by VanderLinden(1993).. whereas their amplitude outside the range plotted in Figure 5. is negligible."," Figure \ref{fig:eigen1} focuses on the region where the eigenfunctions show significant variations, i.e., the conductive layer described by \citet{vanderlinden93}, whereas their amplitude outside the range plotted in Figure \ref{fig:eigen1} is negligible."
" When Figure 5. is repeated for the linear and Gaussian temperature profiles, we find that the eigenfunctions are shifted toward smaller r/R for the linear profile, and to larger r/R for the Gaussian profile."," When Figure \ref{fig:eigen1} is repeated for the linear and Gaussian temperature profiles, we find that the eigenfunctions are shifted toward smaller $r/R$ for the linear profile, and to larger $r/R$ for the Gaussian profile."
" However, the form of the perturbations is very similar to those displayed in"," However, the form of the perturbations is very similar to those displayed in"
should. not je a problem.,should not be a problem.
 The SNe Ia in he current. nearby. rise-time sample are all esseutially at zero redsift. and so if tie observatious were ou the staudard photometric system there would esseutially be uo SED cepeidence of this p'ocess.," The SNe Ia in the current nearby rise-time sample are all essentially at zero redshift, and so if the observations were on the standard photometric system there would essentially be no SED dependence of this process."
 Unfortunately this is uot the case: a signilicaut portiou of tle nearby data cones from unuiltered. observaious., Unfortunately this is not the case; a significant portion of the nearby data comes from unfiltered observations.
 We are not iu a position to use the same techuique to estimate the systematic error ou the nearby data., We are not in a position to use the same technique to estimate the systematic error on the nearby data.
 R99 discussed some of the uucertaiities associated with trausformiug to the standard pass bauds. but cloose o include these effects as large statisical errors rather than as au overall systematic ellect.," R99 discussed some of the uncertainties associated with transforming to the standard pass bands, but choose to include these effects as large statistical errors rather than as an overall systematic effect."
 Therefore. we simply have to trust that our derived statistical errors incorporate systematic uncertainties due to {lor the nearby sample.," Therefore, we simply have to trust that our derived statistical errors incorporate systematic uncertainties due to for the nearby sample."
 We have measurect the rise time [ror ια sample of 73 high redshift (2=0.15—0.9) specroscopically coufirmed Sve [a discovered and observed by the SNLS., We have measured the rise time from a sample of 73 high redshift $z=0.15-0.9$ ) spectroscopically confirmed SNe Ia discovered and observed by the SNLS.
 This determination is roughly 6 imes more srecise than those previously available in this redshift range (AISNOO. COT).," This determination is roughly 6 times more precise than those previously available in this redshift range (AKN00, G01)."
 Our meast'ement for liis sample MNT.y1=19.100.158yay(stat)30.2(syst) days.," Our measurement for this sample is $\tr = 19.10^{+0.18}_{-0.17} \left(\mbox{stat}\right) 
\pm 0.2 \left(\mbox{syst}\right)$ days."
 Using the same analysis techuique on a sample of eight near!w SNe Ia (2« 0.1). we derive a value of 7=19.580.uo days. where the quoted. error jncorporates both satistical aul systematic errors.," Using the same analysis technique on a sample of eight nearby SNe Ia $z < 0.1$ ), we derive a value of $\tr = 19.58^{+0.22}_{-0.19}$ days, where the quoted error incorporates both statistical and systematic errors."
 These differ at the 1.16 level., These differ at the 1.4 level.
 Iu other words. using a οςyusiclerably more precise comparison made possible by ast »stautially )ette| data set. we ind uo coripelling evideuce for auy difference betwee the rise times of nearby and¢instant SNe Ia. It 15 impo‘tant to uuderstaud the limitations of this neasurement in terms of its Constraints on 11eoretical 1iodels.," In other words, using a considerably more precise comparison made possible by a substantially better data set, we find no compelling evidence for any difference between the rise times of nearby and distant SNe Ia. It is important to understand the limitations of this measurement in terms of its constraints on theoretical models."
 As was the case in R99. AINNOO. and GOL. the uncertainties με...ued above are the error iu the mean stretcli-corrected rise times of he two samples. the scatter of rise Limes between individual SNe Ia. However. testing for differences between the WO salupes Is still a very useful €jeck against evoutiouary effects that may e allectiug cosmological analyses uslue HMSNe Τα. The above result suggests ha we are curreuly limited by the systematic uucertainty associated wit1 ]lu perlorming this comparison.," As was the case in R99, AKN00, and G01, the uncertainties presented above are the error in the mean stretch-corrected rise times of the two samples, the scatter of rise times between individual SNe Ia. However, testing for differences between the two samples is still a very useful check against evolutionary effects that may be affecting cosmological analyses using SNe Ia. The above result suggests that we are currently limited by the systematic uncertainty associated with in performing this comparison."
 Tierelore. sigulficant advances will likely require better coustraiuts on the early-time SEDs of SNe Ia. Alternativev. for a large enough sample. it may be possible to Cconstraln {le rise-time by only considering reshifts for whieh the observer and rest-Draine filters inatcl1 paricularly well. minimizing the eerrors.," Therefore, significant advances will likely require better constraints on the early-time SEDs of SNe Ia. Alternatively, for a large enough sample, it may be possible to constrain the rise-time by only considering redshifts for which the observer and rest-frame filters match particularly well, minimizing the errors."
 The qiadratie rise-time model. motivatecl by simple johysical arguments. provides a good fit to the data.," The quadratic rise-time model, motivated by simple physical arguments, provides a good fit to the data."
" Dropping this assumption. we find η=1.850.2 for fx/"" at early times."," Dropping this assumption, we find $n = 1.8 \pm 0.2$ for $f \propto t^n$ at early times."
 Fitting for this extra parameter substantially weakens our coustraiuts onτε. but does uot indicate any discrepancies.," Fitting for this extra parameter substantially weakens our constraints on, but does not indicate any discrepancies."
 Redoiug the fits with a fixed at this value also does not siguilicautly allect any. of our results. simply decreasing all of the measurements of," Redoing the fits with $n$ fixed at this value also does not significantly affect any of our results, simply decreasing all of the measurements of"
disk edge by e.a/to(R)=(P/O? (their Equation 1s).,"disk edge by $c_{s,d}/v_\phi(R) = (\Gamma/\xi)^{1/3}$ (their Equation 18)."
 Combining these two results. the expected Mach umber of the aceretion-driven turbulence is Since fragmentation is avoided only for disks with & of no more than a few. we can take ©~1. aud it iuuediatelv follows that the expected Mach ΙΙ M~1.," Combining these two results, the expected Mach number of the accretion-driven turbulence is Since fragmentation is avoided only for disks with $\xi$ of no more than a few, we can take $\xi \sim 1$, and it immediately follows that the expected Mach number $\mathcal{M} \sim 1$."
 We can apply a similar analysis to real protostellar disks: the Mach uuuber of the turbulence im these disks should follow equation (55))., We can apply a similar analysis to real protostellar disks: the Mach number of the turbulence in these disks should follow equation \ref{machdisk}) ).
" This menus that disks accretiug with ©~1. correspouding to AM~10D AL. vy+ for typical out disk temiperatures To~50 Wy, should have disks whose turbulent velocity dispersions are roughlv trausonic."," This means that disks accreting with $\xi \sim 1$, corresponding to $\dot{M}_{\rm ext} \sim 10^{-5}$ $\msun$ $^{-1}$ for typical out disk temperatures $T\sim 50$ K, should have disks whose turbulent velocity dispersions are roughly transonic."
 This state should prevail during the majority of the main accretion phase., This state should prevail during the majority of the main accretion phase.
 Once tle main accretion pliase euds and the accretion rate drops. the turbulent velocity dispersion should drop to subsouic values.," Once the main accretion phase ends and the accretion rate drops, the turbulent velocity dispersion should drop to subsonic values."
 This represents another prediction from our analysis: class O and class I protostars should have disks with transsonic turbulent velocity dispersions. while class IL aud class IIT sources should have subsonic turbulent velocity dispersions.," This represents another prediction from our analysis: class 0 and class I protostars should have disks with transsonic turbulent velocity dispersions, while class II and class III sources should have subsonic turbulent velocity dispersions."
 As ALMA comes online in the next few vears and provides resolved molecular Luc maps of protostellar disks at a variety of stages in their evolution (o.e.2). we will be able to test this prediction.," As ALMA comes online in the next few years and provides resolved molecular line maps of protostellar disks at a variety of stages in their evolution \citep[e.g.][]{krumholz07d}, we will be able to test this prediction."
 The central approximation we make in our mnocdel is that transport of mass. angular momentum. and energv produced by eravitationally-diiven turbulence can be represented with a local viscous stress tensor.," The central approximation we make in our model is that transport of mass, angular momentum, and energy produced by gravitationally-driven turbulence can be represented with a local viscous stress tensor."
 The validity of this approximation has been the subject of ereat debate in the past decade., The validity of this approximation has been the subject of great debate in the past decade.
 ? show that selt-eravitating disks caunot iu ecneral be modeled with a viscous formali. but that such an approximation may be reasonable for disks near Q=1. the couditiou that we adopt throughout this work. aud that appears to apply to the ealactic and protostellar disks we are mterested iu studving.," \citet{balbus99a} show that self-gravitating disks cannot in general be modeled with a viscous formalism, but that such an approximation may be reasonable for disks near $Q=1$, the condition that we adopt throughout this work, and that appears to apply to the galactic and protostellar disks we are interested in studying."
 Based on a combination of analvtic arguments and local simulations. ? argues that a local prescription is applicable to Q=1 disks that are sufficicutly thin. 5=0.12. aud more recent global smulatious (2???) generally support this result. ο," Based on a combination of analytic arguments and local simulations, \citet{gammie01a} argues that a local prescription is applicable to $Q=1$ disks that are sufficiently thin, $s \la 0.12$, and more recent global simulations \citep{lodato04a, lodato05a, boley06a, cossins09a} generally support this result. \citeauthor{gammie01a}'"
 ον condition for a local transport approxinatiou to apply is woelbsatisBed for ealactie disks at redshifts τν=2 5.1)) and for non-fraeimeutiug protostellar disks .5.1))., 's condition for a local transport approximation to apply is well-satisfied for galactic disks at redshifts $z \la 2$ \ref{sec:diskevol}) ) and for non-fragmenting protostellar disks \ref{sec:protostellar}) ).
 It is mareinally violated for the observed disks at :—2 5.3)). sugeestine that our model should be considered with some caution for them.," It is marginally violated for the observed disks at $z\sim 2$ \ref{sec:highz}) ), suggesting that our model should be considered with some caution for them."
 At a mini the thickness of these disks likely produces different fragimcutatiou behavior than a standard thin disk aualvsis would sugecst (?).., At a minimum the thickness of these disks likely produces different fragmentation behavior than a standard thin disk analysis would suggest \citep{begelman09a}.
 Iu this paper we derive the basic evolution equations for a disk of gas and stars kept iu a state of mareinal eravitational instability by a combination of external accretion. inward migration of eas. aud decay of turbulent motions due to radiative shocks.," In this paper we derive the basic evolution equations for a disk of gas and stars kept in a state of marginal gravitational instability by a combination of external accretion, inward migration of gas, and decay of turbulent motions due to radiative shocks."
 In such a disk. we use the equatious of conservation of mass. angular momentum. and cucrey to derive au equation (22)) that characterizes the iustautaneous rates of mass aud aueular 1iomientuni transport required to maiutaiu the state of mareinal stabilitv. aud we show that this equation has au analytic steady-state solution iu which the disk velocity dispersion (Equation 31)). surface density (Equation 35)). aud rates of transport (Equations 36 and 37)) through the disk are determined by the rate of external infall onto the disk aud the eas mass fraction within it.," In such a disk, we use the equations of conservation of mass, angular momentum, and energy to derive an equation \ref{torquenondim}) ) that characterizes the instantaneous rates of mass and angular momentum transport required to maintain the state of marginal stability, and we show that this equation has an analytic steady-state solution in which the disk velocity dispersion (Equation \ref{sigmaeq}) ), surface density (Equation \ref{coldeneq}) ), and rates of transport (Equations \ref{vreq} and \ref{alphaeq}) ) through the disk are determined by the rate of external infall onto the disk and the gas mass fraction within it."
 We show that disks converge to this steady state ou timescales of order the orbital time. much less than the time over which either the rotation curve or the gas mass fraction changes significautly.," We show that disks converge to this steady state on timescales of order the orbital time, much less than the time over which either the rotation curve or the gas mass fraction changes significantly."
" Based on our analytic solution for the properties of a eravitational instability-domunated disk aud their dependence ou the eas ass fraction and the iufall rate, we are able to eain new insight iuto several processes,"," Based on our analytic solution for the properties of a gravitational instability-dominated disk and their dependence on the gas mass fraction and the infall rate, we are able to gain new insight into several processes."
 We show that the velocity dispersions of both the outer II disks of present dav galaxies aud the main disks of redshift ~2 galaxies can be uuderstood naturally if they are ina state of eravitational instabilitv-reeulated equilibrium., We show that the velocity dispersions of both the outer H disks of present day galaxies and the main disks of redshift $\sim 2$ galaxies can be understood naturally if they are in a state of gravitational instability-regulated equilibrium.
 Moreover. we can uuderstaud the general progression of galactic disks from low values of rotation speed to velocity dispersion ratio. Vinay/7. at hieh redshift to uch higher values today.," Moreover, we can understand the general progression of galactic disks from low values of rotation speed to velocity dispersion ratio, $V_{\rm max}/\sigma$, at high redshift to much higher values today."
 This progression ijs driven primarily by a falloff in galaxy accretion rates and secondarily by the development of disks with stellar velocity dispersion much lager than the eas velocity dispersion. reducing the importance of stars in setting the eravitational instability condition.," This progression is driven primarily by a falloff in galaxy accretion rates and secondarily by the development of disks with stellar velocity dispersion much lager than the gas velocity dispersion, reducing the importance of stars in setting the gravitational instability condition."
 We also predict that the observed rauge of Vyasfo values secu at 2~ is primarily a sequence in eas nies fraction., We also predict that the observed range of $V_{\rm max}/\sigma$ values seen at $z\sim 2$ is primarily a sequence in gas mass fraction.
 Finally. we use the same model to study the velocity dispersions of protostellar disks.," Finally, we use the same model to study the velocity dispersions of protostellar disks."
 We show that our results are iu good agreement with numerical simulations of eravitational instability im disks. aud we predict that velocity dispersions should be traussonic iu class 0 aud I protostars. dropping to subsonic for class IT and IIT SOITCOS.," We show that our results are in good agreement with numerical simulations of gravitational instability in disks, and we predict that velocity dispersions should be transsonic in class 0 and I protostars, dropping to subsonic for class II and III sources."
" Although our attention in this paper is focused on cases that can be solved analytically or nearly so. we close by pointing out that our model. as a result of its eroundiue iu the basic equations of fluid αναος, is also amenable to a more general ununuerical treatment."," Although our attention in this paper is focused on cases that can be solved analytically or nearly so, we close by pointing out that our model, as a result of its grounding in the basic equations of fluid dynamics, is also amenable to a more general numerical treatment."
 One can casily relax our assumptions of coustaut gas fraction. ieelieible influence from stellar feedback. aud a fixed relationship between gas and star velocity dispersion.," One can easily relax our assumptions of constant gas fraction, negligible influence from stellar feedback, and a fixed relationship between gas and star velocity dispersion."
 The resulting equations are identical to the ones we ave already solved. except that they would ποσα to be solved umuercallv.," The resulting equations are identical to the ones we have already solved, except that they would need to be solved numerically."
" There is no fundamental barrier to doing so however. aud the result would be a new method or smnulatius the evolution of mareinally unstable star-forming disks that is intermediate between purely analytic models such as those we have pursued here aud ""ll uniuerieal simulations that can be extremely costly."," There is no fundamental barrier to doing so however, and the result would be a new method for simulating the evolution of marginally unstable star-forming disks that is intermediate between purely analytic models such as those we have pursued here and full numerical simulations that can be extremely costly."
 We plan to pursue this avenue in future work., We plan to pursue this avenue in future work.
" We thauk DDekel. CCaraud. R. IEKlesseu. D. N. LLu. aud MMhuri-Clav for helpful conversations. BBertin. DDokel. EForbes. ντατου, LLodato. ADMIcNally. and RRice for cohunents on the manuscript. and the anouviious referee"," We thank Dekel, Garaud, R. Klessen, D. N. Lin, and Murray-Clay for helpful conversations, Bertin, Dekel, Forbes, Kratter, Lodato, McNally, and Rice for comments on the manuscript, and the anonymous referee"
 , 
"There are two other O2-type stars (Sk—68? 137 and to the north and northwest of 30 Dor, which were BI2253)proposed as runaways by Walborn et al. (","There are two other O2-type stars $-$ $^\circ$ 137 and 253) to the north and northwest of 30 Dor, which were proposed as runaways by Walborn et al. ("
2002).,2002).
" BI2253 has also been observed within the Tarantula Survey, but from preliminary inspection of its spectra, its radial velocity appears consistent with this part of the LMC."," 253 has also been observed within the Tarantula Survey, but from preliminary inspection of its spectra, its radial velocity appears consistent with this part of the LMC."
" The case of #0016 is also reminiscent of the discovery of N11-026 (O2.5 III(f*)), suggested as a runaway by Evans et al. ("," The case of 016 is also reminiscent of the discovery of N11-026 (O2.5 III(f*)), suggested as a runaway by Evans et al. ("
2006).,2006).
" Runaway stars are thought to result from either dynamical interaction in massive dense clusters, or via a kick from a supernova explosion in a binary system, with the more massive star exploding first (see, e.g., Gvaramadze, Gualandris Portegies Zwart 2009, and references therein)."," Runaway stars are thought to result from either dynamical interaction in massive dense clusters, or via a kick from a supernova explosion in a binary system, with the more massive star exploding first (see, e.g., Gvaramadze, Gualandris Portegies Zwart 2009, and references therein)."
 It is generally accepted that R136 is sufficiently young MMyr) that its most massive stars have yet to explode as supernovae., It is generally accepted that R136 is sufficiently young Myr) that its most massive stars have yet to explode as supernovae.
" This implies that, if from R136, #0016 must have been ejected through dynamical interaction, one of the clearest cases to date in support of this mechanism."," This implies that, if from R136, 016 must have been ejected through dynamical interaction, one of the clearest cases to date in support of this mechanism."
" This is vitally important as dynamical interactions in massive clusters are thought to be a possible mechanism for producing stellar mergers and very massive stars (with masses 7 Mo) which might subsequently end their lives as pair-instability supernovae, of broader relevance to the early ages of the universe when such massive stars are thought to be common (e.g., Heger et al."," This is vitally important as dynamical interactions in massive clusters are thought to be a possible mechanism for producing stellar mergers and very massive stars (with masses $>$ $M_\odot$ ) which might subsequently end their lives as pair-instability supernovae, of broader relevance to the early ages of the universe when such massive stars are thought to be common (e.g., Heger et al."
 2003; Gal-Yam et al., 2003; Gal-Yam et al.
 2009)., 2009).
" We thank Alex de Koter, Paco Najarro, and the referee, Hans Zinnecker, for their helpful"," We thank Alex de Koter, Paco Najarro, and the referee, Hans Zinnecker, for their helpful"
1 (FAS) and + (0/15). with a read. noise of 2:36...,"$^{-1}$ (f/8) and $^{-1}$ (f/15), with a read noise of $^{\rm -}$."
 The wavelength range is cispersed quacdraticallvy over the Ποιά at rresolution with blue wavelengths to the north., The wavelength range is dispersed quadratically over the field at resolution with blue wavelengths to the north.
 At both foci. we observed several stellar Dux standards (Lable 1): at δι we also observed. four planetary nebulae.," At both foci, we observed several stellar flux standards (Table 1); at f/8, we also observed four planetary nebulae."
 Observations were also mace in the direction of the South Galactic Pole (Blancd-Lawthorn. Freeman Quinn 1997. hereafter BEQ). and twilight Lats were taken on all nights with and without he etalon.," Observations were also made in the direction of the South Galactic Pole (Bland-Hawthorn, Freeman Quinn 1997, hereafter BFQ), and twilight flats were taken on all nights with and without the etalon."
 We begin by establishing the optical axis of the incomplete calibration rings to better than 0.1 pix., We begin by establishing the optical axis of the incomplete calibration rings to better than 0.1 pix.
 TFhis was achieved with orthogonal distance regression., This was achieved with orthogonal distance regression.
 Next. we azimuthallv bin the calibration rings for all nights to ensure that (a) there were no opto-mechanical shifts. (b) the instrumental response was constant. and (c) the etalon gap zZero-point Was constant from night to night.," Next, we azimuthally bin the calibration rings for all nights to ensure that (a) there were no opto-mechanical shifts, (b) the instrumental response was constant, and (c) the etalon gap zero-point was constant from night to night."
 In (b). the ctalon was found to behave reliably except that there was a slow drift in the optical gap just after twilight on the second. night.," In (b), the etalon was found to behave reliably except that there was a slow drift in the optical gap just after twilight on the second night."
 Three exposures for the Smith LE field were not used., Three exposures for the Smith II field were not used.
 In total. we obtained a total of 2.0 hrs on Smith I. 3.7 hrs on Smith 11. and 2.0 hrs for the olfset sky. position.," In total, we obtained a total of 2.0 hrs on Smith I, 3.7 hrs on Smith II, and 2.0 hrs for the offset sky position."
 We also have about six hours of observation for cach night towards the South Galactic Pole (BEC)., We also have about six hours of observation for each night towards the South Galactic Pole (BFQ).
 Some calibration spectra showed baseline. variations alter binning over dillerent parts of the field., Some calibration spectra showed baseline variations after binning over different parts of the field.
 This arises from —:Iumination (vignette) ellects. stray light. chip structure and CCD fringeing which can be divided out reliably with atfields.," This arises from illumination (vignette) effects, stray light, chip structure and CCD fringeing which can be divided out reliably with flatfields."
 The spectral bandpass seen by individual pixels is roughly wwhere the bandpass centroid declines by from the optical axis to the field edge., The spectral bandpass seen by individual pixels is roughly where the bandpass centroid declines by from the optical axis to the field edge.
 Twilight Hats were ound to be the most reliable except that the Fraunhofer μα»ectrum. leaves its imprint in the data., Twilight flats were found to be the most reliable except that the Fraunhofer spectrum leaves its imprint in the data.
 We divide out 16 solar spectrum. from the fatfield by establishing the mean spectrum and then generating a polar image with this spectrum., We divide out the solar spectrum from the flatfield by establishing the mean spectrum and then generating a polar image with this spectrum.
 Occasionally. the flatfields leave residual fringe structure in the data.," Occasionally, the flatfields leave residual fringe structure in the data."
 Lt is possible to remove this with azimuthal smoothing but potentially informative. intensity variations in the data will be washed out.," It is possible to remove this with azimuthal smoothing but potentially informative, intensity variations in the data will be washed out."
 CCD fringing constitutes the main svstematic error in dilfuse detection ancl provides a major challenge to achieving deeper detection limits than that quoted by ((1994: hereafter D'PVS)., CCD fringing constitutes the main systematic error in diffuse detection and provides a major challenge to achieving deeper detection limits than that quoted by (1994; hereafter BTVS).
 The response of the blocking filter is removed in the twilight division., The response of the blocking filter is removed in the twilight division.
 To obtain discrete. spectra. from the summed. CCD, To obtain discrete spectra from the summed CCD
Ssurvey” (LVIIIS: Ixoribalski e al.,Survey' (LVHIS; Koribalski et al.
 2005)., 2008).
 Since no PROB distance is currently available for NGC 1512. we use lis Local Croup velocity. —— l.tocomputealubble distance of 79.5 Alpe. ," Since no TRGB distance is currently available for NGC 1512, we use its Local Group velocity, =, to compute a Hubble distance of $\sim$ 9.5 Mpc. —"
LVLUS is a large tha aims to provide detailed distributions. velocity. fiels and star formation rates [or a complete sample of nearby. gas-rich galaxies.," LVHIS is a large that aims to provide detailed distributions, velocity fields and star formation rates for a complete sample of nearby, gas-rich galaxies."
 With the Australia Telescope Compact Array (APCA). we observed all LV. galaxies that were detected in the PParkes All-Skv Survey (LIIPASS: Barnes ct al.," With the Australia Telescope Compact Array (ATCA), we observed all LV galaxies that were detected in the Parkes All-Sky Survey (HIPASS; Barnes et al."
 2001. Ixoribalski ct a.," 2001, Koribalski et al."
 2004) and reside south of approx., 2004) and reside south of approx.
 cadeclination., declination.
 The closes neighbours to the NGC 1512/1510 system are (1) the edge-on spiral galaxy NGC 1495 (LILPASS 035844). (2) the ealaxy pair NGC 1487. (LILPASS 0355.42) and (3) the σαaxv ESO0249-C026 (LILPASS 0354.43). all," The closest neighbours to the NGC 1512/1510 system are (1) the edge-on spiral galaxy NGC 1495 (HIPASS J0358--44), (2) the galaxy pair NGC 1487 (HIPASS J0355–42) and (3) the galaxy ESO249-G026 (HIPASS J0354–43), all"
‘To date. more than 40 multiple planetary system have been detected: bevond. our solar svstem.,"To date, more than 40 multiple planetary system have been detected beyond our solar system."
 The commoensurability of orbital periods is. very ubiquitous in. the extrasolar planetary systems., The commensurability of orbital periods is very ubiquitous in the extrasolar planetary systems.
 At present. four resonant pair of planets (GJ ST6. LD 82943. LD 128311. LD 73526) are reported to be trapped in 2:1 mean motion resonance (Macyetetal. 2006).," At present, four resonant pair of planets (GJ 876, HD 82943, HD 128311, HD 73526) are reported to be trapped in 2:1 mean motion resonance \citep{mar01, may04, vog05, tin06}."
. In. recent. vears. numerous researchers. have extensively investigated the dynamics ancl origin of the 2:1 resonance in the planetary systems (CGozdziewski&/Ma-atziset.al. 2009)..," In recent years, numerous researchers have extensively investigated the dynamics and origin of the 2:1 resonance in the planetary systems \citep{Goz01, lee02,
Haj02, Haj06, Haj03, Ji03a, Ji03b, lee04, Kle04, Bea03, Bea06,
Gay08, lee09, voy09}."
" The two resonance variables for 2:1 resonance. Bj=A,PA»|wy. fo=AL2ÀS|xe (where A. zc are the mean longitude and the longitude of periapse respectively. the subscripts 1. 2 denote the inner and outer planets). are categorized to librate about: (190 in svmametric configuration. (2) both 0 and. 180 respectively. in the so-called: antisvmametrie configuration. and (3) other degrees dillerent from 0 or 180 in asymmetric configuration (Iladjidemetriou2002:Jietal.2003a.b.c:)cnugéetal.2003.2006:Lee 2004)."," The two resonance variables for 2:1 resonance, $\theta_1=\lambda_1-2 \lambda_2+\varpi_1$, $\theta_2=\lambda_1-2
\lambda_2+\varpi_2$ (where $\lambda$, $\varpi$ are the mean longitude and the longitude of periapse respectively, the subscripts 1, 2 denote the inner and outer planets), are categorized to librate about: $^\circ$ in symmetric configuration, (2) both $^\circ$ and $^\circ$ respectively, in the so-called antisymmetric configuration, and (3) other degrees different from $^\circ$ or $^\circ$ in asymmetric configuration \citep{Haj02, Ji03a, Ji03b,
Ji03c, Bea03, Bea06, lee04}."
. Phe GJ STG was revealed to be the first 2:1 resonant svstem (Marcy2001) near an M cbwarf star., The GJ 876 was revealed to be the first 2:1 resonant system \citep{mar01} near an M dwarf star.
" The GJ 876 svstem is in apsiclal corotation where the mean motion resonance variables. 6, and 6» librate about Q with quite slight amplitudes."," The GJ 876 system is in apsidal corotation where the mean motion resonance variables, $\theta_1$ and $\theta_2$ librate about $0^\circ$ with quite slight amplitudes."
 On the origin of mean motion resonances in the system. a formation scenario is that they were assembled by migration of planets.," On the origin of mean motion resonances in the system, a formation scenario is that they were assembled by migration of planets."
 In the formation of giant planets. if two planets are massive enough to open gaps and not far away from each other in the disk. the material between the region of them can be rapidly cleared olf.," In the formation of giant planets, if two planets are massive enough to open gaps and not far away from each other in the disk, the material between the region of them can be rapidly cleared off."
 And then. the dissipation of the stull outside the outer planet and inside the inner planet may still force two planets approach cach other.," And then, the dissipation of the stuff outside the outer planet and inside the inner planet may still force two planets approach each other."
 Anyprocess that makes two bodies approach cach other. which are originally separated appropriately. could result in mean motion orbital resonance (Lee&Peale2002:Kleyctal.2005:Massetet 2006)..," Anyprocess that makes two bodies approach each other, which are originally separated appropriately, could result in mean motion orbital resonance \citep{lee02, Kle05,
mas06}. ."
 Lee and collaborators explored the origin and diversity, Lee and collaborators explored the origin and diversity
range in which conditions are stably excellente. or good (PWV 3 mim) are similar (within uncertainties) at Mauna Ixea and ORAL,range in which conditions are stably excellente or good $\leq$ 3 mm) are similar (within uncertainties) at Mauna Kea and ORM.
 This comparison demostrates that excellent. conditions or LR astronomical observations. in terms of percentage and stability. are also possible in sites at relative low altitude (over 2000 m above sea level).," This comparison demostrates that excellent conditions for IR astronomical observations, in terms of percentage and stability, are also possible in sites at relative low altitude (over 2000 m above sea level)."
 This result might be related o the role of the troposphere thickness instead of the site altitude in the LU quality of a particular astronomical site (Garcta-Lorenzoetal.2004).., This result might be related to the role of the troposphere thickness instead of the site altitude in the IR quality of a particular astronomical site \citep{2004SPIE.5572..384G}.
 Indeed. the seasonal behaviour ofthe PWY at ORAL follows a quite similar behaviour than he troposphere thickness at this site (see figure 5 in Lorenzoetal. (2004))).," Indeed, the seasonal behaviour of the PWV at ORM follows a quite similar behaviour than the troposphere thickness at this site (see figure 5 in \cite{2004SPIE.5572..384G}) )."
 The comparison of PWY statistica results between astronomical sites is very complex due lo the variety of techniques and procedure normally used., The comparison of PWV statistical results between astronomical sites is very complex due to the variety of techniques and procedure normally used.
 Dillerent techniques provide dillerent lempor:d coverages and samplings., Different techniques provide different temporal coverages and samplings.
 Aloreover. sonie of the techniques used. to evaluate the water vapor content are sronglv allected by rain and clouds (e.g. radiosondes anc WVIU.," Moreover, some of the techniques used to evaluate the water vapor content are strongly affected by rain and clouds (e.g. radiosondes and WVR)."
 Others are only applicable near the zenith (e.g. racjiometers). whereas GPS can provide continuous PWY esimates depite the elfects of rain. dust or clouds.," Others are only applicable near the zenith (e.g. radiometers), whereas GPS can provide continuous PWV estimates depite the effects of rain, dust or clouds."
 In spite of such dilferences. we summarize in table 3 statistical DWV results for dillerent astronomical sites located at cüllerent. :dtitudes above sea level.," In spite of such differences, we summarize in table \ref{PWV_summarize} statistical PWV results for different astronomical sites located at different altitudes above sea level."
 In the case of Mauna. Wea. we have found. median," In the case of Mauna Kea, we have found median"
The formation of dark matter halos plays a central vole in (he studies of galaxy formation as well as of the large-scale structure of the universe.,The formation of dark matter halos plays a central role in the studies of galaxy formation as well as of the large-scale structure of the universe.
 A widely used analytical theory for the formation of halos is the extended Press-Schechter (PS thereafter) formalism 1991:Bower 1991).. which can be used to model the mass function of dark halos 1914).. the bias parameter as a function of halo mass AIW96).. as well as the formation histories of dark halos (Lacey&Cole1993.1994)..," A widely used analytical theory for the formation of halos is the extended Press-Schechter (PS thereafter) formalism \citep{bond91,bower91}, which can be used to model the mass function of dark halos \citep{ps74}, the bias parameter as a function of halo mass \citep[][hereafter MW96]{mw96}, as well as the formation histories of dark halos \citep{LC1993,LC1994}. ."
 In the simplest excursion-set model of Bondetal.(1991)... the physical properties of dark halos are expected to depend only on halo mass. but not on large-scale environments.," In the simplest excursion-set model of \citet{bond91}, the physical properties of dark halos are expected to depend only on halo mass, but not on large-scale environments."
" However. it has been known for almost a decade (hat neither (he PS mass function nor the MW96 bias model matches well N-bodwv results for halo masses A/<M,. where A, is defined such that the rms linear density fIuctuation within a sphere of mass AL, is ὃς=1.68 (Jing1998:Lee&ShandarinSheth&Lemson1999:Porcianietal.JingTorme"," However, it has been known for almost a decade that neither the PS mass function nor the MW96 bias model matches well $N$ -body results for halo masses $M < M_\ast$, where $M_\ast$ is defined such that the rms linear density fluctuation within a sphere of mass $M_\ast$ is $\delta_c=1.68$ \citep{jing98,lee98,sl99,porciani99,j991,st99}."
n 1999).. (2001) proposed an ellipsoid collapse model to replace the spherical collapse model in the PS theory. and found a better agreement between the theory ancl N-body simulations in both the mass function and the bias function.," \citet{smt01} proposed an ellipsoid collapse model to replace the spherical collapse model in the PS theory, and found a better agreement between the theory and $N$ -body simulations in both the mass function and the bias function."
 Earlier investigations with cosmological N-body simulations did not find any strong environmental dependence of halo properties. such as the spin parameter. concentration and formation time (Lemson&Ixauffinann1999:Percivalοἱal. 2003).," Earlier investigations with cosmological $N$ -body simulations did not find any strong environmental dependence of halo properties, such as the spin parameter, concentration and formation time \citep{lemson99,percival03}."
. ILowever. (he poor resolutions and small dynamical ranges covered in these simulations make it clifficult to detect any signals that are relatively weak or outside the dvnanmical range. and so a signilicant age-dependence of halo clustering cannot be ruled owl.," However, the poor resolutions and small dynamical ranges covered in these simulations make it difficult to detect any signals that are relatively weak or outside the dynamical range, and so a significant age-dependence of halo clustering cannot be ruled out."
 Note also that there were earlier attempts to develop empirical models for the of halo clustering (Taruva& )..," Note also that there were earlier attempts to develop empirical models for the age-dependence of halo clustering \citep{taruya00,Hamana01,YJTS01}."
 With the use ofa very large V-body simulation. Gaoetal.(2005) recently demonstrated convincingly that the clustering strength of dark halos does depend on halo formation time.," With the use of a very large $N$ -body simulation, \citet{gao05} recently demonstrated convincingly that the clustering strength of dark halos does depend on halo formation time."
" In particular. they found that this dependence is strong only for halos with M«AZ, and becomes verv weak lor Af>AZ."," In particular, they found that this dependence is strong only for halos with $M\ll M_\ast$ and becomes very weak for $M>M_\ast$."
" This also explains why Percivaletal.(2003) could. nol detect an age dependence. because (hey focused on halos with Af>M, at high redshifts."," This also explains why \citet{percival03} could not detect an age dependence, because they focused on halos with $M>M_\ast$ at high redshifts."
 Waneοἱal.(2006) have examined (the physical process that may be responsible for the age dependence of halo clustering., \citet{wang06} have examined the physical process that may be responsible for the age dependence of halo clustering.
" They found that halos embedded in clense environments accrete mass less efficiently. (han the spherical collapse model predicts. because the matter to be accreted is ""heated. by the large-scale structure (like(hepancakeheatingconsideredin.e.g..Moetal.2005:Lee 2006)."," They found that halos embedded in dense environments accrete mass less efficiently than the spherical collapse model predicts, because the matter to be accreted is `heated' by the large-scale structure \citep[like the pancake heating considered in, 
e.g.,][]{mo05,lee06}."
. This explains why the old population of simall halos has a higher bias than the voung population of the samemass., This explains why the old population of small halos has a higher bias than the young population of the samemass.
 It also qualitatively explains why the PS mass function and bias functions deviate [rom N-hocly simulations at small masses., It also qualitatively explains why the PS mass function and bias functions deviate from $N$ -body simulations at small masses.
"Poluw rug galaxies (ΓΠΣ) and related objects are peculiar systems that help to understand the formation of galaxies in geuceral. since they represeut extreme cases, xovidiugC» ches ou formation scenarios.","Polar ring galaxies (PRGs) and related objects are peculiar systems that help to understand the formation of galaxies in general, since they represent extreme cases, providing clues on formation scenarios."
" The cearest cases of polar rings. named “kinematically confined"" objects in the Polar Ring Catalog (PRC) of Whitmore et al. ("," The clearest cases of polar rings, named ”kinematically confirmed” objects in the Polar Ring Catalog (PRC) of Whitmore et al. ("
1990). are sable objects frozen iu à peculiar norpholoey. with matter rotating in two ucarly perpendicular planes.,"1990), are stable objects frozen in a peculiar morphology, with matter rotating in two nearly perpendicular planes."
 Nunernieal simulations have shown that this kiud of system could be explained either by simple eas accretion (Beshetuikovy Sotuikova 1997. aud Bournaucd Combes 2003. hereafter DOCU). or by mergers of perpendicularly orieuted disk galaxies (Dekki 1997 1998: BC'O3).," Numerical simulations have shown that this kind of system could be explained either by simple gas accretion (Reshetnikov Sotnikova 1997, and Bournaud Combes 2003, hereafter BC03), or by mergers of perpendicularly oriented disk galaxies (Bekki 1997 1998; BC03)."
 ILlowewer. in the PRC iuost objects do not show such a xhuple appearance with an edec-on central object and an edge-on polar ring. but rather more complex morphologies. iu which they are frozen di a quasi-οπήπι state. avoiding complete relaxation towards a fully mixed svuuuetric svsten," However, in the PRC most objects do not show such a simple appearance with an edge-on central object and an edge-on polar ring, but rather more complex morphologies, in which they are frozen in a quasi-equilibrium state, avoiding complete relaxation towards a fully mixed symmetric system."
" These systems with conrplex morphologies. called polar ring related objects. xovide a unique chance to study PRG formation scenarios in objects in which the initial coupoucuts have not vet clisappeared. if they are the results of merecrs. it may still ο possible to recognize the progenitors in the uurclaxed renunauts. and recoustruct their formation events,"," These systems with complex morphologies, called polar ring related objects, provide a unique chance to study PRG formation scenarios in objects in which the initial components have not yet disappeared – if they are the results of mergers, it may still be possible to recognize the progenitors in the unrelaxed remnants, and reconstruct their formation events."
 This is 10 longer possible in more relaxed iiereer remnants. that rave become elliptical galaxies with ouly faut shells aud ripples.," This is no longer possible in more relaxed merger remnants, that have become elliptical galaxies with only faint shells and ripples."
 Tn this article we preseut the results of photometric observations aud au Hine spectrum of ESO [21-626 (PRC ο in Whitmore et al., In this article we present the results of photometric observations and an line spectrum of ESO 474-G26 (PRC C-3 in Whitmore et al.
 1990. and AAT 0011-213 in Arp Aladore 1987). a unique ealaxy with two almost orthogonal large-scale optical vines (Fig.," 1990, and AM 0044-243 in Arp Madore 1987), a unique galaxy with two almost orthogonal large-scale optical rings (Fig."
 1)., 1).
 ESO 171-026 is in the list of most huninous galaxies (Cappi et al., ESO 474-G26 is in the list of most luminous galaxies (Cappi et al.
 1998). and a strong source of far-infrared (NED!)) aud CO emission (Galletta ct al.," 1998), and a strong source of far-infrared ) and CO emission (Galletta et al."
 1997)., 1997).
 Its unclear spectrum. is «Xf an intermediate type LINERΤΙ (Sekiguchi Wolscnerott 1993)., Its nuclear spectrum is of an intermediate type – LINER/HII (Sekiguchi Wolstencroft 1993).
 The PRC shows optical ciuission-line rofation curves along three position augles., The PRC shows optical emission-line rotation curves along three position angles.
 Observations are detailed iu Sect., Observations are detailed in Sect.
 2. and results iu Sect.," 2, and results in Sect."
 3., 3.
 A anmuuerieal simulation is presented iu Sect., A numerical simulation is presented in Sect.
 4 to illustrate a possible formation mechanism., 4 to illustrate a possible formation mechanism.
 The photometric observatious were obtained in August. 2002. on the 1.6 telescope of the Observatórrio do Pico dos Dias (operated by he MCT/Laboratórrio Nacional de Astroffssica.Drazil). equipped with direct imaging camera #11 and a thick. back-illuniuated 90152015. 13.5-;14 square pixels CCD detector #998. aud scale 07118/pixcl.," The photometric observations were obtained in August, 2002, on the 1.6-m telescope of the Observatórrio do Pico dos Dias (operated by the MCT/Laboratórrio Nacional de Astrofíssica,Brazil), equipped with direct imaging camera 1 and a thick, back-illuminated 2048x2048, $\mu$ m square pixels CCD detector 98, and scale 18/pixel."
 The readout roise was 2.le and the gain. 2.5¢ ‘ADU.," The readout noise was $e^{-}$ and the gain, $e^{-}$ /ADU."
 The data were acquired with standard Joliusou 5. V and Cousins A. 7 banc filters.," The data were acquired with standard Johnson $B$, $V$ and Cousins $R$, $I$ band filters."
 Photometric calibration was made using standard stars from the Landolt (1983) and Graham (1982) lists., Photometric calibration was made using standard stars from the Landolt (1983) and Graham (1982) lists.
 The secing during the observatious was 1733., The seeing during the observations was 3.
 A loe of observations is eiven in Table 1. where Z isthe zenith augle aud Shymeg the extinction corrected sky briehtuess (iu mag 7).," A log of observations is given in Table 1, where $Z$ is the zenith angle and $Sky mag$ the extinction corrected sky brightness (in mag $^{-2}$ )."
 Reduction of the CCD data was performed in the standard manner using the, Reduction of the CCD data was performed in the standard manner using the
Faraday thickness detectable to the present observations is 830 rad m~?.,Faraday thickness detectable to the present observations is 830 rad $^{-2}$.
" Some sources, such as the Radio Arc &Morris1987,1988),, have an RM that changes(Yusef-Zadeh by more than 830 rad m7, so parts of the GC region may be Faraday thick to our observations."," Some sources, such as the Radio Arc \citep{y87,y88}, have an RM that changes by more than 830 rad $^{-2}$, so parts of the GC region may be Faraday thick to our observations."
 Figure 4 shows histograms of ((number of independent spatial beams per degree of A0)) from the entire survey and a smaller region., Figure \ref{histcomp} shows histograms of (number of independent spatial beams per degree of ) from the entire survey and a smaller region.
" Intrinsically, we expect the hhistogram to have contributions from many distinct regions of varying peak aand width."," Intrinsically, we expect the histogram to have contributions from many distinct regions of varying peak and width."
 We found that a Lorentzian profile fits these heterogeneous distributions better than a single Gaussian., We found that a Lorentzian profile fits these heterogeneous distributions better than a single Gaussian.
" For smaller regions, where hhas a single-valued, noise-like distribution, the Lorentzian can also approximate a single Gaussian."," For smaller regions, where has a single-valued, noise-like distribution, the Lorentzian can also approximate a single Gaussian."
" In the limit of a single-valued, noise-like ddistribution, the Gaussian noise is equivalent to a Poisson distribution in the limit."," In the limit of a single-valued, noise-like distribution, the Gaussian noise is equivalent to a Poisson distribution in the limit."
" We use this similarity to approximate the bbin count errors as oy=1+VN4-0.75, where is the number of independent beams in a bin (Gehrels 1986)."," We use this similarity to approximate the bin count errors as $\sigma_N = 1 + \sqrt{N + 0.75}$, where is the number of independent beams in a bin \citep{g86}."
". In §3.2,, we show that comparing RM measured histogram methods to previous work shows that the errors are conservative."," In \ref{poln_comparison}, we show that comparing RM measured histogram methods to previous work shows that the errors are conservative."
" Aside from theoretical expectations, the distribution of vvalues shows that the apparent ccan generally be reliably converted to RM."," Aside from theoretical expectations, the distribution of values shows that the apparent can generally be reliably converted to RM."
" Most values of mmeasured have small offsets from 0°,, with ~50% within +10°,, ~75% within +20°,, and ~90% within +45°.."," Most values of measured have small offsets from , with $\sim$ within $\pm$, $\sim$ within $\pm$ , and $\sim$ within $\pm$."
 This is consistent with the ~9° rrotation expected for RM»+2000 rad πιό., This is consistent with the $\sim$ rotation expected for $\approx\pm2000$ rad $^{-2}$.
" The typical angle change is <1 rrad, so relative angles may be treated as roughly linearly distributed."," The typical angle change is $\ll1$ rad, so relative angles may be treated as roughly linearly distributed."
" The best constraint on the mean hhas a typical error of about1°,, or 220 rad τα”."," The best constraint on the mean has a typical error of about, or $220$ rad $^{-2}$."
" For a change of bbetween our two bands, the Faraday rotation from A—0 cem is49?."," For a change of between our two bands, the Faraday rotation from $\lambda=0$ cm is."
". This typical uncertainty in mmakes the calculation of the intrinsic polarization angle highly uncertain, so no such results are presented here."," This typical uncertainty in makes the calculation of the intrinsic polarization angle highly uncertain, so no such results are presented here."
" Visualizing the images of aand RM is difficult, since the per-pixel sensitivity is poor and varies across the field of view."," Visualizing the images of and RM is difficult, since the per-pixel sensitivity is poor and varies across the field of view."
 Convolution and other image processing techniques can be used to extract this information even in poorly-calibrated VLA data (Rudnick&Brown2009)., Convolution and other image processing techniques can be used to extract this information even in poorly-calibrated VLA data \citep{r09}.
. We tested two statistical techniques to spatially smooth the RM: averaging and histogram fitting., We tested two statistical techniques to spatially smooth the RM: averaging and histogram fitting.
 These methods and a comparison of their results are described in Appendix A.., These methods and a comparison of their results are described in Appendix \ref{meandt}.
" In general, the two methods have similar results."," In general, the two methods have similar results."
" The histogram-fitting method is less sensitive to outliers and has more conservative errors, so it is used in all results described below."," The histogram-fitting method is less sensitive to outliers and has more conservative errors, so it is used in all results described below."
" Since this work is applying arelatively new technique to a complex region, it is important to test the results against known sources."," Since this work is applying arelatively new technique to a complex region, it is important to test the results against known sources."
 This section compares our results shown in Figures 1 and 3 to the RM for specific regions studied previously1997)., This section compares our results shown in Figures \ref{poln_polc} and \ref{poln_padi} to the RM for specific regions studied previously.
. LaRosaetal.(2001) present images of 6 cm polarized intensity near the nonthermal radio filament G359.85+0.39 from VLA data with similar sensitivity and resolution as the present study., \citet{l01} present images of 6 cm polarized intensity near the nonthermal radio filament G359.85+0.39 from VLA data with similar sensitivity and resolution as the present study.
" The two surveys have similar brightness distributions and structure in the polarized emission, particularly the depolarized regions on the southeast and northeast sides of G359.85+0.39 (seealsoLawetal.2008a)."," The two surveys have similar brightness distributions and structure in the polarized emission, particularly the depolarized regions on the southeast and northeast sides of G359.85+0.39 \citep[see also][]{gcl_vla}."
. The similarity shows that the calibration and imaging quality is similar to that of LaRosaetal., The similarity shows that the calibration and imaging quality is similar to that of \citet{l01}.
" Hayneset(2001)..al.(1992) and Tsuboietal.(1986) conducted independent, single-dish surveys near 3 cm, covering a few square degrees of the GC region."," \citet{h92} and \citet{t86} conducted independent, single-dish surveys near 3 cm, covering a few square degrees of the GC region."
" Although depolarization is weaker near 3 cmandtheir beam is larger, there is general agreement between our Figure 1 and their polarized intensity maps."," Although depolarization is weaker near 3 cmandtheir beam is larger, there is general agreement between our Figure \ref{poln_polc} and their polarized intensity maps."
 Figure 5 shows a comparison ofour RMswith the four-band measurements of Tsuboietal. , Figure \ref{rmcomp} shows a comparison ofour RMswith the four-band measurements of \citet{t86}. .
"Near (0°17,0°22) aand (0°1, 0°35), Tsuboietal.(1986) (1986)..find the RM has a maximum of +1000 rad m~?, while the present survey"," Near $(0\ddeg17,0\ddeg22)$ and $(0\ddeg1,0\ddeg35)$ , \citet{t86} find the RM has a maximum of +1000 rad $^{-2}$ , while the present survey"
comprehensive. set of OGLE data.,"comprehensive, set of OGLE data."
 Phase zero has been set to JD2450008.6 corresponding to the peak of the optical flux., Phase zero has been set to JD2450008.6 corresponding to the peak of the optical flux.
 Data from the RATE all sky monitor experiment were obtained from the public archive for RN JO544.17100., Data from the RXTE all sky monitor experiment were obtained from the public archive for RX J0544.1–7100.
 Since the signal from this object is very weak. the data shown in the lower panel of Figure 1 have been averaged over Lo week intervals.," Since the signal from this object is very weak, the data shown in the lower panel of Figure 1 have been averaged over 1 week intervals."
 Looking carefully at the X-ray lishteurve in Figure 1 it is clear that a significant X-ray. signal is only present in the first half of the data run., Looking carefully at the X-ray lightcurve in Figure 1 it is clear that a significant X-ray signal is only present in the first half of the data run.
 Therefore only X-ray data covering the epoch ‘LD 100-1000. were searched for periodic signals., Therefore only X-ray data covering the epoch TJD 100-1000 were searched for periodic signals.
 The raw daily data for this period were then analysed using the same Lomb-Scarele algorithmic technique as the optical data., The raw daily data for this period were then analysed using the same Lomb-Scargle algorithmic technique as the optical data.
 No periodic behaviour was identified in these X-ray data., No periodic behaviour was identified in these X-ray data.
" The peak ASAT X-ray luminosity may be determined from Figure 1 as reaching values of ~0.5 cts/s. ""This corresponds to a peak luminosities of ~6 x Lo’ erg/s and is in good agreement with the BeppoSAX value of S x 107” erg/s quoted earlier in this paper.", The peak ASM X-ray luminosity may be determined from Figure 1 as reaching values of $\sim$ 0.5 cts/s. This corresponds to a peak luminosities of $\sim$ 6 x $10^{35}$ erg/s and is in good agreement with the BeppoSAX value of 8 x $10^{35}$ erg/s quoted earlier in this paper.
 The sources were observed. by the SAO. LOm telescope in 1996. 1999 and 2000.," The sources were observed by the SAAO 1.0m telescope in 1996, 1999 and 2000."
 The exact dates of the individual observations are given in Tables 1 and 2., The exact dates of the individual observations are given in Tables 1 and 2.
 The data were collected using the Teks CCD giving ai field. of approximately 3 arcmin ancl a pixel scale of 0.3 aresee per pixel., The data were collected using the Tek8 CCD giving a field of approximately 3 arcmin and a pixel scale of 0.3 arcsec per pixel.
 Observations were made through standard Johnson UBVRI and Strommeren-Crawlord filters., Observations were made through standard Johnson UBVRI and Strömmgren-Crawford $\beta$ filters.
 The data were reduced. using LRA and Starlink software. and. the instrumental magnitudes were corrected. to the standard system using I region standards.," The data were reduced using IRAF and Starlink software, and the instrumental magnitudes were corrected to the standard system using E region standards."
 In addition. Ht data on RA J0544.1-7100 were obtained from the public archives of the 2ALASS survey of the LMC.," In addition, IR data on RX J0544.1-7100 were obtained from the public archives of the 2MASS survey of the LMC."
 Though the cata were taken about a vear before our optical observations. they are included. here so that possible H1 emission from the circumstellar disk could be investigated.," Though the data were taken about a year before our optical observations, they are included here so that possible IR emission from the circumstellar disk could be investigated."
 There are no reported LR data on RA J0520.5-6932 in the same catalogue., There are no reported IR data on RX J0520.5-6932 in the same catalogue.
 All the resulting photometric magnitudes arc given in ‘Tables 1 and 2., All the resulting photometric magnitudes are given in Tables 1 and 2.
 The average OGLE values over all their data come [rom Ucalski (private communication)., The average OGLE values over all their data come from Udalski (private communication).
 Of course. it is obvious from Figures | and 2 that the sources have not been constant and so we should not expect perfect agreement between the OGLE average values and the specific values reported here.," Of course, it is obvious from Figures 1 and 2 that the sources have not been constant and so we should not expect perfect agreement between the OGLE average values and the specific values reported here."
 A red spectrum of RA 05441-7100. is shown in Figure 5. was obtained from the 1.98]. SAAQO observatory on 9 January 1999 using the Cassegrain spectrograph with the SELO2 CCD detector.," A red spectrum of RX J0544.1-7100 is shown in Figure 5, was obtained from the 1.9m SAAO observatory on 9 January 1999 using the Cassegrain spectrograph with the SITe2 CCD detector."
 Though of a low signal-to-noise ratio. the spectrum. clearly shows Ha in emission. though little can be determined about the line shape.," Though of a low signal-to-noise ratio, the spectrum, clearly shows $\alpha$ in emission, though little can be determined about the line shape."
 The La line has an equivalent width of EW = -TXLX and a central position of 65684LA., The $\alpha$ line has an equivalent width of EW = $\pm$ and a central position of $\pm$.
.. Blue spectroscopy of both sources were obtained on Ist November L999 using the ESO 1.52-m telescope at La Silla Observatory. Chile.," Blue spectroscopy of both sources were obtained on 1st November 1999 using the ESO 1.52-m telescope at La Silla Observatory, Chile."
 “Phe telescope was equipped. with the Boller Chivens spectrograph | #332 holographic, The telescope was equipped with the Boller Chivens spectrograph + 32 holographic
t(R) =,t(R) =
t(R) =T,t(R) =
t(R) =Ti,t(R) =
t(R) =Tia,t(R) =
diszuption radius we estimate that the progenitor core density would be py~6«LOPALpc7.,disruption radius we estimate that the progenitor core density would be $\rho_0 \sim 6 \times 10^5 M_\odot {\rm pc}^{-3}$.
 The width of he peaks comprising the disk (aud heuce the progenitor core radius] is estimated to be 21| pe., The width of the peaks comprising the disk (and hence the progenitor core radius) is estimated to be 2–4 pc.
 However. NOC LIsGA is much more distant than M3L. and consequently he eccentric disk mass is probably much higher than that of M31. probably exceeding LOSAL.. and this is too massive or a elobular cluster.," However, NGC 4486A is much more distant than M31, and consequently the eccentric disk mass is probably much higher than that of M31, probably exceeding $10^8 M_\odot$ and this is too massive for a globular cluster."
 However as in the case of M31s estimated progenitor. the central densities. core radius aud otal mass are reasonable for a galaxy core.," However as in the case of M31's estimated progenitor, the central densities, core radius and total mass are reasonable for a galaxy core."
 Iu equation(2)) we have assumed an isothermal profile for the stellar density profile of the background. galaxy., In \ref{rt_def}) ) we have assumed an isothermal profile for the stellar density profile of the background galaxy.
 ILlowever. as shown by multiple studies (e.g... Faberetal.1997:Laueretal. 1998)) galaxy cores at small radii are seldom ft with a smfacebrightuess profile proportional to Ro? (where R is the observed angular radius from he nucleus ou the skv) corresponding to an isothermal deusity profile X+7.," However, as shown by multiple studies (e.g., \citealt{faber,lauer98}) ) galaxy cores at small radii are seldom fit with a surfacebrightness profile proportional to $R^{-1}$ (where $R$ is the observed angular radius from the nucleus on the sky) corresponding to an isothermal density profile $\propto r^{-2}$."
 huages from the IST have been wart of a major effort to classify the nuclear stellar profiles in carly-type galaxies. resulting in the classification of ight profiles iuto two categories. galaxies with shallow Inner cusps. denoted ‘core-type profiles aud galaxies with “power-law light profiles (Laucretal.1995:Faberct L997).," Images from the HST have been part of a major effort to classify the nuclear stellar profiles in early-type galaxies, resulting in the classification of light profiles into two categories, galaxies with shallow inner cusps, denoted `core-type' profiles and galaxies with `power-law' light profiles \citep{lauer95,faber}."
.. Power-law type ealaxy cores tend to have the steepest nuclear surface brightuess profiles p()xR with 5 nearly equal to 1 in only the the most extreme cases., Power-law type galaxy cores tend to have the steepest nuclear surface brightness profiles $\mu(R) \propto R^{-\gamma}$ with $\gamma$ nearly equal to 1 in only the the most extreme cases.
 Most of the galaxies studied by Faberctal.(1997) had shallower profiles., Most of the galaxies studied by \citet{faber} had shallower profiles.
 This nuplies that equatiou(3)) somewhat underestimates the transition radius ia most ealaxies. particularly in the galaxies classified as core-tvpe iu which 5 is close to zero.," This implies that \ref{r_t}) ) somewhat underestimates the transition radius in most galaxies, particularly in the galaxies classified as core-type in which $\gamma$ is close to zero."
 We can regard our estimated radius (equation 3)) as a lower limit ou the transition radius for all but the galaxies with the steepest nuclear profiles., We can regard our estimated radius (equation \ref{r_t}) ) as a lower limit on the transition radius for all but the galaxies with the steepest nuclear profiles.
 We now consider the problem of two mereie galaxies. both with more complex. stellar density profiles aud both with massive black holes.," We now consider the problem of two merging galaxies, both with more complex stellar density profiles and both with massive black holes."
 We can approximately describe the primary as having a transition radius eiven by equation(3)) Gvhlich is actually a lower limit as mentioned above)., We can approximately describe the primary as having a transition radius given by \ref{r_t}) ) (which is actually a lower limit as mentioned above).
 If the secondary has a power-law form for its density profile then it will not completely disrupt at a particular radius. like the nou-siugular isothermal sphere or Wine models.," If the secondary has a power-law form for its density profile then it will not completely disrupt at a particular radius, like the non-singular isothermal sphere or King models."
 One component of the Nuker surface brightuess profile (Faberctal.L997) is the break radius. ríj where the slope of the surface brightuess profile changes.," One component of the Nuker surface brightness profile \citep{faber} is the break radius, $r_b$, where the slope of the surface brightness profile changes."
 For profiles denoted core-type. within rj. the surface brightuess (aud so deusitv) rises much less steeply with decreasing radius than outside it.," For profiles denoted core-type, within $r_b$, the surface brightness (and so density) rises much less steeply with decreasing radius than outside it."
 Because the surface brightuess for core- galaxies onlv iucreases slightly within r4. we can associate the break radius aud deusitv estimated at this radius. with the core radius ry and density py which we have used in the previous sections.," Because the surface brightness for core-type galaxies only increases slightly within $r_b$, we can associate the break radius and density estimated at this radius, with the core radius $r_0$ and density $\rho_0$ which we have used in the previous sections."
 This lets us determine if the bulee or stellar component ofthe secondary galaxy will survive tidal truncation to within the transition radius of the primary., This lets us determine if the bulge or stellar component of the secondary galaxy will survive tidal truncation to within the transition radius of the primary.
 A secondary with break radius excecdine its own transition radius sj2ry (likely to he true for huninous galaxies) will almost completely disrupt when the tidal truucation radius reaches its break radius., A secondary with break radius exceeding its own transition radius $r_b > r_t$ (likely to be true for luminous galaxies) will almost completely disrupt when the tidal truncation radius reaches its break radius.
 The black hole of the secondary will then spiral in to the uucleus of the primary. leaving the stellar compoucut of the secondary behind.," The black hole of the secondary will then spiral in to the nucleus of the primary, leaving the stellar component of the secondary behind."
 ILlowewer for bulges and fainter cllipticals. we expect ry<r.," However for bulges and fainter ellipticals, we expect $r_b < r_t$."
 For example. M28 profile is steeply rising well within its transition radius s;c9O.5pc and dmereases in density to its last measured poiut (Laueretal.1998).," For example, M32's profile is steeply rising well within its transition radius $r_t \sim 0.5$ pc and increases in density to its last measured point \citep{lauer98}."
. Iu this case the core would be continuously disrupted. though we expect a change iu the surface brightuess profile of the resulting disk set bv the location at which the tidal tyuncatiou radius of the secondary was equal to its own transition radius.," In this case the core would be continuously disrupted, though we expect a change in the surface brightness profile of the resulting disk set by the location at which the tidal truncation radius of the secondary was equal to its own transition radius."
 The continuous stripping of the secondary iu the presence of the secondarys black hole may prevent the formation of a lopsided disk., The continuous stripping of the secondary in the presence of the secondary's black hole may prevent the formation of a lopsided disk.
 The radius of tidal disruption can be estimated by comparing the mean deusitv of the object to that of the ohject disuptiue it (e.e.. Sridhar&Tremaine1992)).," The radius of tidal disruption can be estimated by comparing the mean density of the object to that of the object disrupting it (e.g., \citealt{sridhar}) )."
 At a distance rz from the nucleus of the primary galaxy. the mean cdeusity resulting from the black hole is p(r)=Mandia where Moj4 is the mass of the priuuuvs black hole.," At a distance $r$ from the nucleus of the primary galaxy, the mean density resulting from the black hole is $\bar \rho(r) = M_{bh,1}/({4 \pi r^3 \over 3})$ where $M_{bh,1}$ is the mass of the primary's black hole."
 We have asstuned that &—rj where rgj ds the trausition radius of the primary. so the contribution from the prinarys bulee is negligible.," We have assumed that $r<r_{t,1}$ where $r_{t,1}$ is the transition radius of the primary, so the contribution from the primary's bulge is negligible."
 For a secondary galaxy with a power-law law density profile (see Binney&Tremaine1987. section 2.10) the mass integrated out to radius s is so the mean density within this radius for radii s2ry. outsüde the transition radius of the secondary.," For a secondary galaxy with a power-law law density profile (see \citealt{B+T} section 2.1e) the mass integrated out to radius $s$ is so the mean density within this radius for radii $s> r_{t,2}$ outside the transition radius of the secondary."
 We set the mean deusity of the primary within r (due to the primarys massive black hole) equal to the that of secoudary. (within stripping radius s). p(s)=ptr). and find that When the secondary is at a distance + from the nucleus of the primary. it will be stripped out to a distance s frou its own nucleus. where s is related to r by the previous equation.," We set the mean density of the primary within $r$ (due to the primary's massive black hole) equal to the that of secondary (within stripping radius $s$ ), $\bar\rho(s)=\bar\rho(r)$, and find that When the secondary is at a distance $r$ from the nucleus of the primary, it will be stripped out to a distance $s$ from its own nucleus, where $s$ is related to $r$ by the previous equation."
" The stripping radius. s. is the same size as the distauce from the prinaryv's nucleus. kr. at a radius where Ay,4 ds the mass of the primarys black hole."," The stripping radius, $s$, is the same size as the distance from the primary's nucleus, $r$, at a radius where $M_{bh,1}$ is the mass of the primary's black hole."
 When s estimated frou equation(16)) exceeds + then the secondary cueults the nucleus of the primary., When $s$ estimated from \ref{r3}) ) exceeds $r$ then the secondary engulfs the nucleus of the primary.
 For a disrupted secondary galaxy core to produce au eccentric disk. we expect that the radius at which the secondary cheults the nucleus of the primary. rj. should be smaller than the transition radius of the primary (ya).," For a disrupted secondary galaxy core to produce an eccentric disk, we expect that the radius at which the secondary engulfs the nucleus of the primary, $r_e$ , should be smaller than the transition radius of the primary $r_{t,1}$ )."
 Tu this case. the core of the secondary. will survive intact within the trausition of the secondary.," In this case, the core of the secondary will survive intact within the transition of the secondary."
 A reasonable condition for formation of an eccentric disk via the diszuption of a ealaxy with a power-law deusity profile should be, A reasonable condition for formation of an eccentric disk via the disruption of a galaxy with a power-law density profile should be
in estimating its luminosity aud temperature: stars with an cS lave individual reddening estimates based on observed spectral types. while stars labeled P had photometry corrected assuning the modal extinction derived for cluster ienibers by Rebulletal.(2006).,"in estimating its luminosity and temperature; stars with an `S' have individual reddening estimates based on observed spectral types, while stars labeled `P' had photometry corrected assuming the modal extinction derived for cluster members by \citet{Rebull}."
. Finally. the star added to the distance sample using the esin/ measurement cataloged bv Rebulletal.(2006) lacks a correspondiug esiní error estimate: its entrv du Table 3. lists None in the esin/ error colin.," Finally, the star added to the distance sample using the $v \sin i$ measurement cataloged by \citet{Rebull} lacks a corresponding $v \sin i$ error estimate: its entry in Table \ref{tab:sinidata} lists `None' in the $v \sin i$ error column."
 Figure 5 shows the distribution of spectral types for stars in the distance sample., Figure \ref{spectypedist} shows the distribution of spectral types for stars in the distance sample.
 Most of the stars in the distance sample have spectral types in the ranee of I1 to AL., Most of the stars in the distance sample have spectral types in the range of K4 to M3.
 Observational uncertaiuties affect the shape of the measured sin/ distribution., Observational uncertainties affect the shape of the measured $\sin i$ distribution.
 Iu order to obtain a reliable distance estimate. we must account for these uncertainties iu our modeled values of sn.," In order to obtain a reliable distance estimate, we must account for these uncertainties in our modeled values of $\sin
i$ ."
" To do this. we first estimate unicertainties iu P. esiu/. Toy, aud 1 from the observational data: a Monte Carlo simulation is then used to incorporate these error distributions iuto the modeled value of sin/."," To do this, we first estimate uncertainties in $P$, $v \sin i$, $T_{eff}$ and $L$ from the observational data; a Monte Carlo simulation is then used to incorporate these error distributions into the modeled value of $\sin i$."
 The rotation periods of T Tauri stars can be measured with high precision from their light curves., The rotation periods of T Tauri stars can be measured with high precision from their light curves.
 The errors associated with such measurements are usually ou the order of1%., The errors associated with such measurements are usually on the order of.
.. Tn some cases. coufotuding factors such as aliasing iu the helt curves or the presence of multiple starspots can increase these errors dramatically (6.8.Herbstetal.2002).," In some cases, confounding factors such as aliasing in the light curves or the presence of multiple starspots can increase these errors dramatically \citep[e.g.][]{Herbst02}."
. Tlowever. the ummber of cases in which these effects oceur is typically siuall.," However, the number of cases in which these effects occur is typically small."
" Since the period errors are very stall compared to the errors associated with other T Tauri mcasurements (1.6. Iuumosity and 7,.,,). we simply assuue that fractional errors in period are normally distributed with a standard deviation of characteristic of the typical errors in pre-nain sequence stellar period measurements."," Since the period errors are very small compared to the errors associated with other T Tauri measurements (i.e. luminosity and $T_{eff}$ ), we simply assume that fractional errors in period are normally distributed with a standard deviation of, characteristic of the typical errors in pre-main sequence stellar period measurements."
 For the remaining variables it is possible to ascertain sole measure of the actual errors from the observed data., For the remaining variables it is possible to ascertain some measure of the actual errors from the observed data.
 Fractional esins errors for the distance sample were calculated using the relationship between HR aud A(esin’) given in 63..., Fractional $v \sin i$ errors for the distance sample were calculated using the relationship between $R$ and $\delta (v \sin i)$ given in $\S$ \ref{testing_uncertainties}. .
 The resulting error distribution is consistent with a normal distribution with c=20%.., The resulting error distribution is consistent with a normal distribution with $\sigma$.
" Fractional crrors in τε were determined Ὃν comparing values of Try, calculated using individual spectral type-based reddening corrections or simply adopting the medal reddening for all cluster members.", Fractional errors in $T_{eff}$ were determined by comparing values of $_{eff}$ calculated using individual spectral type-based reddening corrections or simply adopting the modal reddening for all cluster members.
 The T.;; οπου «istribution caleulatecdk using this prescrption is consisent with a normal distribution with a =QUA., The $_{eff}$ error distribution calculated using this prescrption is consistent with a normal distribution with $\sigma$ =.
" We expο, however. that calculating Tyy¢¢ values by adopting the clusters modal reddening is less accurate than deriving reddenuiues frou observed spectral types."," We expect, however, that calculating $_{eff}$ values by adopting the cluster's modal reddening is less accurate than deriving reddenings from observed spectral types."
 Indeed. we note a clear relationship between the OT.¢¢ value measured for cach star aud its spectroscopic reddening estimate.," Indeed, we note a clear relationship between the $\delta$ $_{eff}$ value measured for each star and its spectroscopic reddening estimate."
" This suggests that the difference between the two Try, estimates is dominated by the errors iutroduced by adopting the modal reddeniug. and that the resultiug error distribution overestimates the actual errors associated with the Γρ values derived using iudividual spectral tvpe-based. reddening corrections."," This suggests that the difference between the two $_{eff}$ estimates is dominated by the errors introduced by adopting the modal reddening, and that the resulting error distribution overestimates the actual errors associated with the $_{eff}$ values derived using individual spectral type-based reddening corrections."
 Siularh the errors im L are estimated from the difference between Iluminosities. calculated assuuiug a reddening derived from cach stars spectral type aud those caleulated assumniusg an overall reddening for the cluster.," Similarly, the errors in $L$ are estimated from the difference between luminosities calculated assuming a reddening derived from each star's spectral type and those calculated assuming an overall reddening for the cluster."
 We fud that the hunuinositv error distribution inuplied by these distinct £ estimates lies within the bounds of the normaldistribution with a=35% sueecstecd bv Ihutiiuun(2001) ascharacteristic of hnunünositv errors in pre-niadn sequence stars., We find that the luminosity error distribution implied by these distinct $L$ estimates lies within the bounds of the normaldistribution with $\sigma$ suggested by \citet{Hartmann_agespreads} ascharacteristic of luminosity errors in pre-main sequence stars.
 Our technique for modelingthe distribution of sin/ in NGC 2261 borrows heavily from Preibisch&Suüthl (1997)., Our technique for modelingthe distribution of $\sin i$ in NGC 2264 borrows heavily from \citet{Preibisch}. .
". We define the modeled value of sn. (in), as"," We define the modeled value of $\sin i$ , $(\sin i)_{m}$ as"
"Consider a relativistic blast wave (hat propagates near a stellar edge pre-explosion mass density gradient of the form ox2"". where 2=(ft,—r). R, is the stellar radius and r ds (he distance trom the star center.","Consider a relativistic blast wave that propagates near a stellar edge pre-explosion mass density gradient of the form $\rho \propto
z^n$, where $z=(R_*-r)$, $R_*$ is the stellar radius and $r$ is the distance from the star center."
 In. compact stars. where shocks are more likely (o become relativistic. (he envelope is radiative and (vpically vn223. which is the value that we use (throughout the paper.," In compact stars, where shocks are more likely to become relativistic, the envelope is radiative and typically $n \approx 3$, which is the value that we use throughout the paper."
 The shock acceleration stops when the width of the shock is comparable (o 2 ancl the shock breaks out., The shock acceleration stops when the width of the shock is comparable to $z$ and the shock breaks out.
" As we show in appendix A. this takes place in (he shell where the pre-explosion optical depth to the stellar edge is 70.015,. where , is the shock Lorentz factor."," As we show in appendix A, this takes place in the shell where the pre-explosion optical depth to the stellar edge is $\tau \sim 0.01 \g_s$, where $\g_s$ is the shock Lorentz factor."
" Namely. breakout takes place at 2~0.015,/(pig). where Hy020.2cn/g is the Thompson cross section per unit ofmasst."," Namely, breakout takes place at $z \sim 0.01 \g_s/(\rho \kappa_T)$, where $\kappa_T \approx 0.2 {\rm
~cm^2/g}$ is the Thompson cross section per unit of."
. This sets the maximal Lorentz [actor of the shock ancl we refer to that shell as theshell., This sets the maximal Lorentz factor of the shock and we refer to that shell as the.
 Unlike the Newtonian case. the observed emission from relativistic breakout is not dominated by the shell.," Unlike the Newtonian case, the observed emission from relativistic breakout is not dominated by the shell."
 Instead it is dominated by the shell where the pre-explosion optical depth Toeol] (he. zoe l/(psq)) which we therefore refer to asshell aud denote ils properties by the subscript 4.," Instead it is dominated by the shell where the pre-explosion optical depth $\tau \sim 1$ (i.e., $z \sim 1/(\rho
\kappa_T)$ ), which we therefore refer to as and denote its properties by the subscript $_0$ '."
 In case of properties that evolve with (ime. (his subscript refers (ο their value in the shell right after it is crossed by the shock.," In case of properties that evolve with time, this subscript refers to their value in the shell right after it is crossed by the shock."
 Given the above assumptions the system is completely defined by the lollowing parameters of the breakout shell:, Given the above assumptions the system is completely defined by the following parameters of the breakout shell:
D configuration of 2007 sununer. while the A=1.3 uuu continuum was obtained iu the E coufiguratiou of 2007 sununucr.,"D configuration of 2007 summer, while the $\lambda = 1.3$ mm continuum was obtained in the E configuration of 2007 summer."
 Each data set was taken with oue or two side bands of a 500 ΛΠΙ bandwidth in each sinele-side baud for the continuum observations., Each data set was taken with one or two double-side bands of a 500 MHz bandwidth in each single-side band for the continuum observations.
 Two or one extra bands were assigned to à CO rotational transition oor »0)., Two or one extra bands were assigned to a CO rotational transition or ).
 The CO rotational transition data are presented in another paper with other molecular transition data., The CO rotational transition data are presented in another paper with other molecular transition data.
 The details of each observatiou are listed in Table 1.., The details of each observation are listed in Table \ref{tab_obs}.
 Two aud three pointing mosaics have becu done to better cover the larger bipolar outflow regions for the ttransition towards LIlis IRS 3 aud LI1157. respectively. at À=1.3 umm.," Two and three pointing mosaics have been done to better cover the larger bipolar outflow regions for the transition towards L1448 IRS 3 and L1157, respectively, at $\lambda = 1.3$ mm."
 For this study. the northwest pointing data of Lillis IRS 3 aud the central poiutiug data of L1157 were used.," For this study, the northwest pointing data of L1448 IRS 3 and the central pointing data of L1157 were used."
 The Multichaunel. Tage Reconstruction. Tmaee Analysis. and Display (MIRIAD.?) tools have beeu cluploved to reduce and analyze data.," The Multichannel Image Reconstruction, Image Analysis, and Display \citep[MIRIAD,][]{sault1995} tools have been employed to reduce and analyze data."
 In additiou to normal procedures (lineleneth. baudpass. fux. aud eaiu calibrations). shadow-defected data have been flagged iu the E configuration data.," In addition to normal procedures (linelength, bandpass, flux, and gain calibrations), shadow-defected data have been flagged in the E configuration data."
 Shadowing iudicates cases of an antennas lue-ofsight interrupted by other antennas and usually appears in low clevation observations of conrpact configurations., Shadowing indicates cases of an antenna's line-of-sight interrupted by other antennas and usually appears in low elevation observations of compact configurations.
 The normal effects of shadowing are reduction and distortion of incident autenua power and abnormal gain jumps., The normal effects of shadowing are reduction and distortion of incident antenna power and abnormal gain jumps.
 Therefore. to obtain reliable results the shadow-detected data were flaeeed iu the compact E configuration.," Therefore, to obtain reliable results the shadow-defected data were flagged in the compact E configuration."
 Further special atteutiou needs to be given ou fiux calibration for studies involving flux comparison between different waveleueths like this study;, Further special attention needs to be given on flux calibration for studies involving flux comparison between different wavelengths like this study.
 To minimize errors caused by primary flux calibrators. we used the same fle. calibrator (Uranus) at both waveleneths except L1157. which used NNVC319 at aand Maus atnun.," To minimize errors caused by primary flux calibrators, we used the same flux calibrator (Uranus) at both wavelengths except L1157, which used MWC349 at and Mars at."
 We expect and uncertainties of flux calibrations at aand 2.7 wun. respectively. based on the CARMA conuuissioniusg task of flux calibration.," We expect and uncertainties of flux calibrations at and 2.7 mm, respectively, based on the CARMA commissioning task of flux calibration."
 During a colllissiouing period extending to longer than [| mouths. 12 calibrator (quasar) fluxes had been monitored by CARAIA.," During a commissioning period extending to longer than 4 months, 12 calibrator (quasar) fluxes had been monitored by CARMA."
 As à result. the least virving case showed about deviation in flux.," As a result, the least varying case showed about deviation in flux."
 When considerimg the intrimsic variability. of quasars. if is expected that CARMA flux calibrations have about LO154 uncertainties.," When considering the intrinsic variability of quasars, it is expected that CARMA flux calibrations have about $10-15$ uncertainties."
 As a result. we consider and uucertainties at aand 2.7 mua. respectively.," As a result, we consider and uncertainties at and 2.7 mm, respectively."