ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 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."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.," 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."4" 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 $R_{\rmn{h,max}}$."5" 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."6 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."7 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."8 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."9 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."10 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."11" 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."12 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."13 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 3.14 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 $R_{\rmn{wind}} > 3.75 15\times 10^{11} \rmn{cm}$ ) is $\sim 2 \times 10^{-10} \msol \rmn{yr^{-1}}$."16 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 $10^{-7} 17\msol \rmn{yr^{-1}}$ made by \citet{kot} to explain the hydrogen column density inferred from absorption line spectra."18 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."19 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."20 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."21 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.22 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}.23. 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 $\nu \propto \alpha c_{\rmn{s}}^2$."24 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).25 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.26 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."27 This could lead to extremely high accretion rates during a subsequent outburst., This could lead to extremely high accretion rates during a subsequent outburst.28 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."29 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-mass30Our adopted (My. (V-I)) calibration is shown in Figure 10.,"Our adopted $_V$, (V-I)) calibration is shown in Figure 10."31 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 $\sim 2.85$."32" Finally. Figure LL plots the (M,. (I-J)) relation."," Finally, Figure 11 plots the $_I$, (I-J)) relation."33 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 $1.45 < (I-J) < 1.65$, $9.2 < M_I < 11.2$ )."34 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."35 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."36 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."37(Madan ct 11996: Friagaa Torlevich 1998: Steidel 1999).,(Madau et 1996; Friaçaa Terlevich 1998; Steidel 1999).38 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.39 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)."40 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)."41 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)."42 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)."43 lence. a similar mass netallicity relation can be expected for quasars.," Hence, a similar mass – metallicity relation can be expected for quasars."44 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)."45 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 $z \geq 3.5$.46 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."47 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)."48 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.49 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 $Z/Z_\odot \simeq 4 \,{\rm to}\,5$."50 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."51 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 $z\ga 4$."52We 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).53" 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 $_o \simeq 65$ $^{-1}$ $^{-1}$, $\Omega _M = 0.3$, $\Omega _\Lambda = 0.7$ )."54 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.55 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)."56As 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.57 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.58 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.59 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.60 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."61 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."62 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."63 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 $5\fdg5 64\times 5\fdg5$."65 ‘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.66 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)."67 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."68a one-zone model.,a one-zone model.69 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}70 for temperatures from 10 to 250 K, and from \citet{mckee82}71 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."72 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."73 The average ratio of the narrow and broad line clumps are used., The average ratio of the narrow and broad line clumps are used.74 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 $_{13}$ ."75 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."76" 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."77 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 $(4.5\pm3.5)\times10^{3}$ $^{-3}$."78 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 $(8.5\pm1.5)\times10^{2}$ $^{-3}$.79 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.80" 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."81 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.82 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 $\rm\Sigma_{SFR}$."83 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."84 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).85 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."86 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."87 The cincunmnuclear gas is in general in solid body rotation., The circumnuclear gas is in general in solid body rotation.88 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.89 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."90 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}91 in their low resolution $^{12}$ CO(J = 1–0) map."92" 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."93 We use the AIPS taskGAL to determine the dynamical center., We use the AIPS task to determine the dynamical center.94 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."95 The deduced kinematic parameters are sununarized in Table 6., The deduced kinematic parameters are summarized in Table 6.96 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.97 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}.98 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."99 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."100 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."101 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},"102 is due to the asvuuuetric intensity distribution., is due to the asymmetric intensity distribution.103 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."104 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}.105. 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."106 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."107 Their orientation is almost perpendicular to the major stellar bar., Their orientation is almost perpendicular to the major stellar bar.108 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.109 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)."110 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}."111. 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}.112. 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."113 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."114 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.115" 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}."116. 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."117 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 1097.118 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."119 The schematic sketch is shown in Figure 17.., The schematic sketch is shown in Figure \ref{fig-model}. .120 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."121 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"122", f top)presentsseveralhubsofine spacedaliasperiods.",", top) presents several hubs of fine-spaced alias periods."123"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}))."124T hehubshortwardso ff 1.5yieldsanellipsoidallightcurveatanorbitalperiodP ~22h., The hub shortwards of $f = 1.5$ yields an ellipsoidal light curve at an orbital period $P \sim 22~\mathrm{h}$.125" 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."126 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.127" 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 $1.907-2.381~\mathrm{cyc/d}$."128" 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 $P_2 = 1290.4695~\mathrm{d}$ yield an acceptable light curve."130" 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 $P_\mathrm{ph} = 0.4695(01)~\mathrm{d}$."131" 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."132" 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."133" 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."134" 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."135" 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."136" 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."137 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.138 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 $\log g = 4.5$.139" 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)."140 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.141" 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)."142"T he""spectroscopic"" periodat firstglanceprovidesa , bottom)."," The ""spectroscopic"" period at first glance provides an acceptable fit to the data \\ref{ec13phrv_fig}, bottom)."143"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."144"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."145Wethere f oredoubtthephysicalrelevanceo , We therefore doubt the physical relevance of this signal.146fti ," Again we point out that the longest photometric data set excludes the ""spectroscopic"" period for the light curve."147"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."148 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).149 Aen1650 of this system clearly e-reso.dO Voution spectroscopy., Further investigation of this system clearly requires time-resolved high-resolution spectroscopy.150" 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."151" 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."152" 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}. ."153". 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} "154reduction 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."155 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."156 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}."157 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.158 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."159 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."160" 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.161 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."162 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."163" 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."164 In the sun-centered simulation spherical coordinate svstem. the origin of," In the sun-centered simulation spherical coordinate system, the origin of"165qsPhe 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."166 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."167 1997. IXnapen. Shlosman Peleticr 2000. Fathi et al.," 1997, Knapen, Shlosman Peletier 2000, Fathi et al."168 2006. ΙΟ et al.," 2006, Riffel et al."169 2008)., 2008).170 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."171 EThese jets: can have a considerable: impact. on this medium., These jets can have a considerable impact on this medium.172 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.173" 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).174 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."175 Because of the observational, Because of the observational176"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.177 1991)., 1994).178" 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.179 1993: Palla et al., 1993; Palla et al.180 1995) Musa]., 1995) ].181 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."182" 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."183 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 $10\;{\rm K}\leq \tg \leq 10^4\;{\rm K}$."184 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."185 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 of186in an accretion disk to explain the (win kllz-QDPOs in LMXDs.,in an accretion disk to explain the twin kHz-QPOs in LMXBs.187 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}.188 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."189 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 $I_* \approx 0.5-1.5 \,190(10\;\textrm{ km})^2 \, M_\odot$."191 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."192 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.193 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."194 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."195 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}. .196straight line to some degree of accuracy.,straight line to some degree of accuracy.197 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 $^{-1}$."198" 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."199 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."200 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.201 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 $^{-1} \pm 6$."202 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 $^{-1}$.203 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."204 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."205 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.206 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)."207 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."208 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."209 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."210 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"211"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."212 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."213" 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."214" 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."215 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."216" 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."217 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 $\sigma_{\rm KDP}\la 1\,\kms$."218 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."219" 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 $\bar v_{KDP}220= 16.4\,\kms$, $A_x = 2.6\,\kms\,\rm deg^{-1}$ , $A_y = 4.9\,\kms\,\rm221deg^{-1}$ , $\sigma_{\rm KDP}=5\,\kms$ , and $\zeta_{\rm KDP}=0.673$."222 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."223 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.224 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 0.225 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."226 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.227 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."228 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."229 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."230 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 4.2.231 In order to make the simulations tractable. we ignore metallicity information.," In order to make the simulations tractable, we ignore metallicity information."232" 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.233 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."234 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."235 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 $2\,\sigma$."236 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."237 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}) )."238" 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."239 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.240 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."241 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}) ).242 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."243 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.244 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.245 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.246 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."247 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."248 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."249 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."250 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 $30\,251\kms$ since close tidal"252and 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 $_2$.253" 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}254 model of protostar AFGL 2591 agree with our observed abundances while their column densities for, and are lower than our observed values."255 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 ."256. 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}) )."257" 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.258. 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."259 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.260 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."261.. 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.262. 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."263" The observed irafio agrees with the branching ratio for the destruction ofCIT,.", The observed ratio agrees with the branching ratio for the destruction of.264. Based on these results «πλοία also be very abunudaut., Based on these results should also be very abundant.265 We now attempt fo construct a physical aud econietrical iiodel ofix., We now attempt to construct a physical and geometrical model of.266 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, lugo04,267debuizer05, kraus06}, we propose a scenario in which the molecular absorption presented here comes from a circumstellar disk."268 We will cxamine other possibilities fist., We will examine other possibilities first.269 Figure 8. depicts the possible scenarios allowed by the available observations., Figure \ref{fig:cartoon} depicts the possible scenarios allowed by the available observations.270 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.271" 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.272 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."273" 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."274 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."275. 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.276 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."277 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}."278. 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."279 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.280 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."281 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."282 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."283 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}) )."284 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.285 The knots may be material stripped fromthe disk., The knots may be material stripped fromthe disk.286 The stellar, The stellar287 high-2 , ${\it z}$ 288Here. 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."289 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."290 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."291. Three initial coucitious are required. of which we lave already identified two:.," Three initial conditions are required, of which we have already identified two:."292. 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."293 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 $g_{-2} (\theta)$."294 We thus have a one-parameter family of outer flow solutionW., We thus have a one-parameter family of outer flow solutions.295 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 $g_{-2} (\theta)$."296 We obtained each solution by asstuning dillerent values of gà(x)., We obtained each solution by assuming different values of $g_{-2} (\pi)$.297 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."298Notice also that all curves attaiu the same value at πο.,Notice also that all curves attain the same value at .299 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 $g_{-2} (\pi)$ ."300orders 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."301" 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 $z = 1100$."302" 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."303 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."304 Phis fact was confirmed by the maximum likelihood analysis., This fact was confirmed by the maximum likelihood analysis.305 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.306 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."307 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."308 The 5 mJv data may be allected bv svstematic surface density gradients., The 5 mJy data may be affected by systematic surface density gradients.309 The results are consistent with an unchanging underlying power spectrum., The results are consistent with an unchanging underlying power spectrum.310 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)."311 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).312 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}) ).313 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 $A_{\ell m} = \sum_i Y_{\ell314m}^*(i)$ are drawn independently from Gaussian distributions such that the normalization satisfies $|A_{\ell m}|^2 = \sigma_0$ ."315" 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 $|A_{\ell m} -316\sigma_0 I_{\ell m}|^2/J_{\ell m}$ (equation \ref{eqclpeeb}) )."317 ligure 8. plots the distribution of observed: values of Adesmodosως , Figure \ref{figalmdist} plots the distribution of observed values of $|A_{\ell m} - \sigma_0 I_{\ell m}|^2/J_{\ell m}$.318We 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."319 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).320 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}.321 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.322" 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."323 Thus equation 14 must be amended such that Gr)xexpμα|2e)o0]., Thus equation \ref{eqalmdist} must be amended such that $P(x) \propto \exp{[-x/(1+2e)\sigma_0]}$.324 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."325 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.326 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.327 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. (3282003).,2003).329 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 $w(\theta) \approx (1 \times 10^{-3}) \,330\theta^{-0.8}$ for angles up to a few degrees."331 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}332 allows us to derive the equivalent $C_\ell$ spectrum if we assume that this power-law extends to all angular scales."333 In Figure 9. we overplot the resulting prediction on the measurements and find an excellent. fit.," In Figure \ref{figwth}334 we overplot the resulting prediction on the measurements and find an excellent fit."335 This is initially surprising: the angular correlation, This is initially surprising: the angular correlation336dynnamics.,namics.337 11904., 1904.338 oobtained., obtained.339 cllusters., lusters.340 staar dynamics, ar dynamics341and nuclear reactions can substantially affect the productivity of p-nuclei.,and nuclear reactions can substantially affect the productivity of p-nuclei.342 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.343" 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."344 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 3rd345columns,][]{Lodd2003}."346" All these sources are associated with core-collapse supernovac,", All these sources are associated with core-collapse supernovae.347Photo-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, Raus2002,348 Haya2008} is currently regarded as the most successful scenario."349" 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”."350 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}.351. 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.352" 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)."353 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)."354" 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 $^{94}$ .355 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{Hoff1996,356 Meye1998b} at $T_9 \sim 4-3$ leads to the formation of these p-nuclei with $N \le 50$."357" 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.358" 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 in359Table~4,][]{Wana2009}."360" 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, $^{96,361 98}$ Ru, $^{102}$ Pd, $^{106, 108}$ Cd, $^{113}$ In, $^{115}$ Sn, and for a heavy p-isotope $^{152}$ Gd."362 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."363 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 $\nu364p$ -process, as follows."365 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}."366 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."367 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 $f >368f_\mathrm{max}/10$ to each p-isotope abundance in Figure 6 (the standard model with $Y_\mathrm{e, 3}$ ranging between 0.5 and 0.7)."369 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 4.370 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 $\sim 0.6$."371 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]”."372 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 $A = 108$.373 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 $\nu374$ p-process can account for the origin of the heavy p-isotopes up to $A375= 152$ as well, if $Y_\mathrm{e, 3} \approx 0.65$ (Figure 6) is achieved in the neutrino-driven outflows."376 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."377 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 $\gtrsim 10^8$.378 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 $f = 10^5$ ).379 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."380 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.381 A word of caution for the ποοσπα isotopes ds needed hore., A word of caution for the molybdenum isotopes is needed here.382" 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 $Y_\mathrm{e, 3} = 0.53-0.54$ ."383 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 5.3.384" 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."385 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}. ."386" 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."387 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)."388" 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."389" 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}. .390 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.391 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.392" 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)."393 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 $Z \ge39456$ ), and the nuclear masses $40 \le Z \le 50$ ) on the $\nu$ p-process pathway."395 Our result, Our result396saturation 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}."397 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."398 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, $\tau$."399" 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}."400 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."401 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."402 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."403 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.404 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)."405 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.406 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."407 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."408 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.409 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."410 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.411 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."412 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.413 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.414 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}) )."415 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.416 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}.417 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}. .418" 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."419 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."420 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."421 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}."422 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}.423 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.424 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}."425 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}.426 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."427 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."428 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.429 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"430"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."431" 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."432" 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)."433" 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."434" 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."435 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.436 Eq. (, Eq. (437"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 $P_{i,ad}(k) = A\,438k^n$ 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)."439" 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).440 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.441" 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. ("442"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."443 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.444 We see the Poisson plateau (Eq., We see the Poisson plateau (Eq.445 8) at large k's which drops with decreasing mass., 8) at large k's which drops with decreasing mass.446 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.447the Milky Was.,the Milky Way.448 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."449 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.450 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."451 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."452 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."453 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."454 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."455 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."456 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."457 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."458 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."459 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."460 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.461 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."462 “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.463 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.464 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."465 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.466 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)."467 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."468 2000)., 2000).469 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."470 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."471 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.472 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)."473" 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."474 Indeed. Odenkirchen ct al. (," Indeed, Odenkirchen et al. ("4752001) 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."476 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.477 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."478 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."479the LIL regions in M. 101 (Ixennicutt. Carnett 1996).,the HII regions in M 101 (Kennicutt Garnett 1996).480 A single line logarithmic gradient is found., A single line logarithmic gradient is found.481 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.482 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 $S_{23}$.483 ‘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.484 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."485 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."486 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."487 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."488 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)."489 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.490 Regarding LIL regions of higher metallicity. the composedparameter SoOo; ," Regarding HII regions of higher metallicity, the composedparameter $S_{23}/O_{23}$ "491"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 $1.4\times10^{10}$ $^{-3}$."492 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."493 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."494 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)."495 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.496 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.)."497 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)."498 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."499 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.500 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."501 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. ("5021) 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.503 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."504 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."505 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. ("5062) 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."507 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. (5083) 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."509" 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."510 It can be expressed as:, It can be expressed as:511cohbunon 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)."512 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).513" 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."514 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."515 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. ("5161991) 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.517" 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."518 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., ("5191999) 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."520 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."521 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)."522" 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.523 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."524 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.525" 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 $dN_e/dE_e \propto E_e526^{-\Gamma}$."527"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 $E_e528= m_e \gamma_{\rm tr}$ where the electron spectrum $dN_e/dE_e$ has the same value for the thermal and nonthermal components."529" 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 $\gamma_{\rm530tr}$ is the lower limit of the Lorentz factor distribution of the nonthermal component, and hence there are no nonthermal electrons at $E_e < m_e \gamma_{\rm tr}$ ."531" 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 $\gamma_{u} =53210^5$, 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 $T_d =533$ 10 eV, following the conventional thermal models (e.g. 1994)."534 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, $\tau_T$ ."535 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 $\alpha_X \approx5361.9$, e.g., 1994; 1997; 1998) when we set $\tau_T =5370.24$."538 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.539 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.540 We then trace the Comptonization process using a Monte Carlo method., We then trace the Comptonization process using a Monte Carlo method.541 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."542 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)."543 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."544 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 U03.545 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."546 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."547" 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."548 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."549 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.550 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."551 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."552 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."553" 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"554"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."555 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."556 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."557 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."558 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"559exited 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."560 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."561". 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."562 Lor radial modes the total energy. dissipation rate Lis η., For radial modes the total energy dissipation rate $D$ is $D_{\rm p}$.563 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."564 “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."565 3., 3.566 In the lower panel of Fig., In the lower panel of Fig.567 6. we plot the quantity A?P.," 6, we plot the quantity $\lambda^2P$."568 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.569 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."570 With the help of equation (21) and the data eiven in Fig., With the help of equation (21) and the data given in Fig.571" 3 we can evaluate amplitudes for radial modes with I,« 0.", 3 we can evaluate amplitudes for radial modes with $D_{\rm p}<0$ .572 In Fig., In Fig.573" 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."574 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."575 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."576 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.577 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."578 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."579 We should stress that the amplitude: estimates in lig., We should stress that the amplitude estimates in Fig.580 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)."581 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"582magnitude.,magnitude.583 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}) ).584 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."585 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."586 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 $\pm$ $^{-1}$."587 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."588" 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."589 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."5905. Using these final clipped. mean velocities. AV = OSZEO. ss!.," Using these final clipped mean velocities, $\Delta V$ = $\pm$ $^{-1}$."591 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) $\times$ $^{-6}$."592 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."593 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."594 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."595 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 1830–211.596 From a colmparisou of the auumouia inversion lines (Ieukel ct al., From a comparison of the ammonia inversion lines (Henkel et al.597 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. ("5982008) fud cousisteucv. at a lo level. of Ap/pe = 1.9<10 ©,"2008) find consistency, at a $\sigma$ level, of $\Delta\mu$ $\mu$ = $\times$ $^{-6}$."599 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.600 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."601. 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."602" 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."603 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."604 1998). (, 1998). (6051) 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.606 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.607 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.608 2008). (, 2008). (6092) 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. (61020082) aud Meuteu et al. (,2008a) and Menten et al. (6112008).,2008).612 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)."613 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."614 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. ("6153) 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."616 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 $\pm$ $^{-1}$."617 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)."618 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.619 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."620 19814: Stutzki Winnewisser 1985a.b).," 1984; Stutzki Winnewisser 1985a,b)."621" 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.622 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."623HESS 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.624 2008)., 2008).625 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).626 ILESS J13534-020 lies towards the southern border of this remnant., HESS J1858+020 lies towards the southern border of this remnant.627 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)."628 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.629 They proposed. using a simplified analytical approach described in Torres et al. (," They proposed, using a simplified analytical approach described in Torres et al. ("6302003). 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."631 In a more recent paper. Paron et al. (," In a more recent paper, Paron et al. ("6322011) 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."633 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."634 Paron et al. (, Paron et al. (6352011) 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 G35.6-0.4.636 Here. we focus on a more in-depth analvsis of this possibility.," Here, we focus on a more in-depth analysis of this possibility."637where 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 $n_{\chi}$."638 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."639 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.640" 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."641 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."642 This is illustrated in figure 1 Ü, This is illustrated in figure \ref{fig:relicabund}~ \cite{Jungman:1995df}.643 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."644 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.645"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 ."646. 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."647 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."648 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.649 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."650 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."651 In other cases. theZ-PEG lits were very poorly constrained.," In other cases, the fits were very poorly constrained."652 Furthermore. there was a tendeney. for the fits to congregate around. z2.," Furthermore, there was a tendency for the fits to congregate around $z\sim2$."653 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."654 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."655 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."656 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.657 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."658 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."659 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."660 The photometric redshifts were therefore not. included. in. this paper., The photometric redshifts were therefore not included in this paper.661 The continuum is well detected (Pig. 2)).," The continuum is well detected (Fig. \ref{continua}) ),"662 but we see no emission or absorptions lines., but we see no emission or absorptions lines.663 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 $z\sim 1$.664 Vhe continuum is well detected (Fig. 2)).," The continuum is well detected (Fig. \ref{continua}) ),"665 but we see no emission or absorption lines., but we see no emission or absorption lines.666 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 $z\sim 1.2$.667 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."668" Phe redshift is based on a single emission ine at Aun,=6051AA.", The redshift is based on a single emission line at $\lambda_{\rm obs}=6051$.669. 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."670. 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."671 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."672 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.673 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."674 No line or continuum emission was detected in the. 3600s FORS2 spectrum., No line or continuum emission was detected in the 3600s FORS2 spectrum.675 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."676 “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."677 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"678AR«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}."679. ~ (Franco-Ieruáudez," $\sim$ \citep{francheretal04,rodrigetal07,galvmadetal08}."680 of spectral type earlier than B3., of spectral type earlier than B3.681 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, $c_i \sim68210$ km/s, they would have lifetimes of roughly $10^4$ yr."683 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."684 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 $\sim68510^5$ yr if this model is correct."686 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}."687. 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 Paper688I]{petersetal10}."689. 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."690 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."691 However. massive stars require accretion af rates excocding LO1 AML. | ," However, massive stars require accretion at rates exceeding $10^{-4}$ $_{\odot}$ $^{-1}$ "692Giant 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.693" 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}."694.. 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.695 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}."696. 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."697 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}."698. 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."699 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?"700 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}?"701? 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}."702. 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."703 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}."704 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?705 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}. ."706The 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.707 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?708 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?709 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}."710. 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."711 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."712 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."713 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."714"(qc 5/3) or Ivaichnan-tvpe(q=3/2). themomentum-dependence becomes a,=1/3 and A,=1/2. respectively,","$q\simeq5/3$ ) or $q=7153/2$ ), themomentum-dependence becomes $\alpha_p = 1/3$ and $\alpha_p=1/2$, respectively."716" 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 $\alpha_p=0$."717" 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}."718. 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}) )719 are In the s-parameter space. (he solution of Eq. (4))," are In the $\gamma$ -parameter space, the solution of Eq. \ref{dif}) )"720" 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.721 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.722" 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 $t_{\rm cool}=1/[\beta_s p]$.723 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."724 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 $\eta \geq 1$."725" 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 $\gamma_c \simeq72610^6~(v_A / 0.01 c)~(1~\mathrm{G}/B)^{1/2}\eta^{-1/2}$."727 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}) ).728" 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 $B_{cr}=m^2c^3/e\hbar$."729 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.730 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 $\nu^{c}_{\rm peak}\propto \gamma_{c}$ .731Note that for an electron distribution of the form of eq. (6)), Note that for an electron distribution of the form of eq. \ref{max}) )732 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 $\propto733\exp[-(\gamma/\gamma_{c})^{\beta}]$ , the corresponding cut-off in the synchrotron spectrum appears much"734Svensson (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."735 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 $\tau$."736 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 $N_{\rm H} = (3.42 \pm 0.3) \times 10^{20}737\rm \, cm^{-2}$, and a second to model absorption within Holmberg II for which wefix the metallicity at $Z = 0.07 Z_{\sun}$ \cite{mirioni02}."738. 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.739 We [ound an adequate fit with yo /Dok = 273.5/285., We found an adequate fit with $\chi^2$ /DoF = 273.5/285.740 The best fit parameters are reported in Table 1.., The best fit parameters are reported in Table \ref{specfits}.741 Phe best fit column density is (3.720.5)107+em.27.," The best fit column density is $(3.7 \pm 0.5) \times 10^{21} \rm \, cm^{-2}$."742 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."743 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}.744.. 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."745 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."746 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).747" 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."748. 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 2003.749 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."750 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."751 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 2002.752 All of the narrow band filters have a bandwidth., All of the narrow band filters have a bandwidth.753 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."754 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."755 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.756 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."757" 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.758 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.759 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 $0.3\arcsec$."760 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}.761.. 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.762 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."763 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."764 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."765 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.766 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.767" 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}. ."768 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"769"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 $\theta_t =770\theta_j$."771 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."772 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.773 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."774 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 $\pm10\arcdeg$."775" 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."776 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."777 KjPu 8., KjPn 8.–778" 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}."779. /; 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}. ."780 AL 1-10., M 1-16.–781 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."782 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.783 2 corresponds to the radius of the small ionized nebula., 2 corresponds to the radius of the small ionized nebula.784 /; for the (wo additional jet components in Fig., $t_j$ for the two additional jet components in Fig.785 2 are 1050 vr and 740 vr (Schwarz1992)., 2 are 1050 yr and 740 yr \citep{sch92}.786. M 2-9., M 2-9.–787 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}."788 AL 1-92., M 1-92.–789 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}790 is from a geometrical method using Doppler shifts in the jets."791" /; 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 $\tan792\theta_t$."793 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."794 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."795 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.796centered al z22withastandarddeviationof1.,centered at $=$ 2 with a standard deviation of 1.797M9 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."798"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: $E_{peak}=E_{peak,obs}(1+z)$."799 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 $S_{bolo}$."800 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."801 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."802 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."803" 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."804 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."805 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."806 2011)., 2011).807 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.808 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."809 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.810 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. ("8112010).,2010).812 This discrepancy is the main topic of this paper., This discrepancy is the main topic of this paper.813" 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."814 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}) )."815 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."816 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.817 1993. Fishinan Meegan 1995).," 1993, Fishman Meegan 1995)."818 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.819 1995)., 1995).820 We then apply a bolometric correction with (a=—1.0 and ~]= —2.0)., We then apply a bolometric correction with $\alpha=-1.0$ and $\beta=-2.0$ ).821 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.822 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."823" 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 $E_{peak,obs}$."824" 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."825 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.826 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"827"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."828 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}."829" 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."830" 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."831 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}832 installed on the Hubble Space Telescope that enables us to study these absorbers at low redshift with unprecedented efficiency."833" 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."834" 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."835" 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."836 The structure of the paper is as follows., The structure of the paper is as follows.837 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.838" 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)."839" 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."840" 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 $\Delta z\sim12$."841 Conclusions are summarized in 8 6., Conclusions are summarized in $\S$ 6.842" 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 $^{-1}$ $^{-1}$ ."843" 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."844" 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 $0.15<z<0.35$ ."845 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 $z\la$ 0.35846Iu 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."847 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."848 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.849 The PSF at £=0 shows the typical instantaneous speckle., The PSF at $t=0$ shows the typical instantaneous speckle.850 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."851 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."852 2007)., 2007).853 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.854" 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."855"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."856 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."857 In general. the PSF isophotes change with radius.," In general, the PSF isophotes change with radius."858 Thus. the measurement somewhat depeuds ou the weighting scheme.," Thus, the measurement somewhat depends on the weighting scheme."859" 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."860" 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 $\sim0.7\arcsec$."861 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 $0.5\tan^{-1}( \epsilon_{\times}/ \epsilon_{+})$."862 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.863 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."864 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.865 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).866 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."867" 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."868 The amplitude of the correlation function decreases with exposure time., The amplitude of the correlation function decreases with exposure time.869 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."870 The 61 c, These results have been validated with a series of 15 sec exposures of the globular cluster NGC2419 on the Subaru telescope.871m 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.872 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.873 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)."874 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."875 The CCD flatness techuologv aud testing is reviewed in Takacs et al. (, The CCD flatness technology and testing is reviewed in Takacs et al. (8762006) and Radcka et al. (,2006) and Radeka et al. (8772009).,2009).878 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.879 Several veudors are now exceeding these flatuess specifications., Several vendors are now exceeding these flatness specifications.880 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}.881 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}.882 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."883 Several features are noteworthy in Figure 8.., Several features are noteworthy in Figure \ref{fig_ellipticity_focal_plane_no_atmos}.884 First. the aberrationdnuduced ellipticity is larec.," First, the aberration-induced ellipticity is large."885 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."886 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."887 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."888 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."889 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."890 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."891 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.892 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"893depending on the time the heat pulse is switched on.,depending on the time the heat pulse is switched on.894" 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."895" 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."896 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.897" 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."898" 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."899" 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."900" 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."901" The peak of the distribution is around 3 MK, as planned."," The peak of the distribution is around 3 MK, as planned."902" 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."903" 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."904 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."905 Fig., Fig.906" 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 $\lambda 3.21$ $log ~ T[K] ~ = ~ 7.4$ )."907" 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}) )."908 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.909 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}) ).910 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."911" 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."912" 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"913Is 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?"914 “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."915 Vhis corresponds to an angular scale of ~ 2microarcseconds., This corresponds to an angular scale of $\sim2$ microarcseconds.916 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."917 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."918 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."919 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."920 Given the ring-like nature of (Jauncy et al., Given the ring-like nature of (Jauncy et al.921 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. ("92219907).,1997).923 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."924 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."925 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)."926 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."927 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."928 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."929 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."930 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."931 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."932 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.933 Hence. such variability could inlluence svstems sullering gravitational microlensing.," Hence, such variability could influence systems suffering gravitational microlensing."934 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."935 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."936 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."937 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."938 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.939 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"940This 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.941 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.942 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."943 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."944 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 $\Delta945m_{15}$, 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}."946. 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."947 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 $\alpha = 0.052$."948 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 $\alpha =9490.123$."950 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."951 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 $\alpha^\prime \Delta m_{15}$.952 This value is consjistent. with 7820.18 obtained by Hamuyetal.(1996)., This value is consistent with $\pm0.18$ obtained by \cite{Hamuy96}.953. 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.954 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 $R_V$ ).955" 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)."956 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}957 and \ref{fig:hostlf_i} with solid histograms for the $r$, $g$, and $i$ -passbands."958 Hostless SNe are indicated by 5iacled istograun at rightimnost bin., Hostless SNe are indicated by shaded histogram at rightmost bin.959 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."960 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."961 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."962" 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"963structures 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}.964. 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.965 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."966 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."967 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."968 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."969 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."970 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.971 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."972 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."973 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.974. 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."975" 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}."976". 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}."977".. 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."978 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."979" 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. $t_{\rm SB} = \max \{\, t_{\rm ff},\, t_{\rm980cool}\,\}$ ."981 Then star formation activitv ⋝↸∖∶↴∙⊾↕∐↴∖↴⋜↧↑⊺∠∖∶⊺∠⇪↖⊺∠⊨↴⊔∙↖↖↽∐↸∖↥⋅↸∖⊺∠⇪↕↴∖↴↑↕∐∖∐⋜↧↕∪ ≯∪↥⋅∐↓⋜↧↑↕∪∐↑↕⊔↸∖↻↥⋅↸∖≺∐≼⊳↑↸∖≺↧↴⋝∙↖⇁↑∐↸∖↴∖↴↑⋜⋯≼↧⋜∐⋅≼↧⊏↕⋟≋ ⊔∪≼∐∖↕∏⇀⋜↧↸⊳↸∖⋅↖⇁∙∖↽≼⊲∪↕↸∖↕∩," Then star formation activity begins at $t_{\rm s} =982t_{\rm F} + t_{\rm SB}$, where $t_{\rm F}$ is the halo formation time predicted by the standard EPS model \citep{LC93}."983⋂∶≩⋝⋈↕↕≯∤∖↕↴∖↴↕⋜∐⋅∶↴∙⊾↸∖↥⋅↑∐⋜⋯ he IInbble nue at the halo redshift. we define he system as aginihalo. o. a starless salaxx.," If $t_{\rm984s}$ is larger than the Hubble time at the halo redshift, we define the system as a, i.e. a starless galaxy."985 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."986" 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 $T_{\rm K} = T_{\rm IGM}987(1+\delta)^{\gamma-1}$ , where $\gamma = 5/3$ is the"988photous 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."989 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."990 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.991 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."992 Fie. 1))., Fig. \ref{xmm_img}) ).993 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.994 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)."995 Nevertheless. there is no couseusus ou its plivsical origin vet.," Nevertheless, there is no consensus on its physical origin yet."996 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."997 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."998 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."999 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."1000 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. ("10012009) 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.1002" 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 $^{-1}$.1003 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.1004 200.b Manchester et al., 2004; Manchester et al.1005 2005)., 2005).1006 To advance the analvsis we preseuted earlier (Tetzlaff et al., To advance the analysis we presented earlier (Tetzlaff et al.1007 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)."1008 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. ("10092011) 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.1010 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. ("10112010) (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.1012 2009 and Tetzlaff et al., 2009 and Tetzlaff et al.1013" 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."1014 1993: Taylor Cordes 1993: Yuan et al., 1993; Taylor Cordes 1993; Yuan et al.1015 2010):, 2010);1016applying 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}."1017 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.1018 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}."1019 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}1020 to suggest that the current best estimate from these two methods is $8\pm1$."1021. 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}."1022 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.1023 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."1024 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.1025 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)."1026 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."1027 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.1028 ? 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."1029 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}."1030 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).1031 Unfortuately that colour regime applies to most of our galaxy sample., Unfortunately that colour regime applies to most of our galaxy sample.1032 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."1033 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}. ."1034 9 estimated the SFR density in four different redshift bins, \cite{Dale2010} estimated the SFR density in four different redshift bins10350.3 auc 0.5 are probably members of the cluster.,0.3 and 0.5 are probably members of the cluster.1036 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.1037 Stars with E(B—V)>0.5 could be background objects., Stars with $E(B-V)>0.5$ could be background objects.1038 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.1039" 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."1040 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).1041 Such large distances seem iniprobable. because they imply unreasonably large absolute magnitudes.," Such large distances seem improbable, because they imply unreasonably large absolute magnitudes."1042 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."1043 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.1044 We examine this possibility further in the uext subsection., We examine this possibility further in the next subsection.1045" 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)."1046 Fig 3 shows these ciagranms., Fig 3 shows these diagrams.1047 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."1048 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.1049 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}.1050 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."1051 These objects may be PMS stars (see section [.)., These objects may be PMS stars (see section 4.).1052 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.1053 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}."1054. 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.1055 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."1056 Radial velocity measuremenW. would test this possibilitv., Radial velocity measurements would test this possibility.1057 The stars with lower reddening occupy au area above the main sequeuce., The stars with lower reddening occupy an area above the main sequence.1058 Most of these stars are [oreerouud objects with different distance moduli., Most of these stars are foreground objects with different distance moduli.1059 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.1060 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.1061 Further observations would place better limits ou its distance., Further observations would place better limits on its distance.1062 The other foreground stars lie just above the main sequence., The other foreground stars lie just above the main sequence.1063 Their position is consistent with the clistance calculated, Their position is consistent with the distance calculated1064Figure 1 shows the average velocity of the H53a line.,Figure 1 shows the average velocity of the $\alpha$ line.1065 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).1066 This velocity pattern is consistent with either rotation or outflow., This velocity pattern is consistent with either rotation or outflow.1067 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.1068" First, the observed velocity gradient is perpendicular to the CO outflow."," First, the observed velocity gradient is perpendicular to the CO outflow."1069 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.1070" 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}."1071". 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 $-30^\circ$ )."1072" 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."1073" 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."1074 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.1075" 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."1076" 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."1077" 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}."1078". 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."1079 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.1080" 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 $."1081" 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}. ."1082 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.1083 Zhang&Ho(1997) measure a velocity gradient in NH3 of 500 kms! pc! , \citet{ZhangHo1997} measure a velocity gradient in $_3$ of 500 $^{-1}$ $^{-1}$ 1084rale 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.1085 90110)., 2011c).1086 Hsieh et al. (, Hsieh et al. (108720110) 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.1088 Fernandez οἱ al. (, Fernandez et al. (10891997) 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."1090 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."1091 The red plate shows only a hint of this tail., The red plate shows only a hint of this tail.1092 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."1093 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."1094 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."1095 2011)., 2011).1096 Dased on a statistical dvnamical model. Bottke et al. (," Based on a statistical dynamical model, Bottke et al. ("10972002) 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.1098 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.1099 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."1100 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."1101 It. need. not be re-described in detail here., It need not be re-described in detail here.1102 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."1103 2011)., 2011).1104 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."1105 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."1106 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"1107It 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)."1108 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 $\xi$."1109 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 $|\theta'(\xi_1)| = 0.0424$."1110 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."1111 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."1112 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."1113 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."1114 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."1115 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."1116 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."1117 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."1118 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."1119 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."1120 We may reverse the above argent aud view eq. (80)), We may reverse the above argument and view eq. \ref{crit1_2}) )1121 as a condition. on Z/M., as a condition on $R/M$.1122 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."1123 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."1124 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.1125 We cun similarly evaluate the second criterion. eq. (73)).," We can similarly evaluate the second criterion, eq. \ref{crit2}) ),"1126 in terms of Lauc-Euidenu functions., in terms of Lane-Emden functions.1127 The angular nonmentum Jtr) of the matter euclosed within radius r.," The angular momentum $J(r)$ of the matter enclosed within radius $r$ ,"1128of 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}."1129. 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 $M_{\rm bol} = +4.74$."1130" 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."1131 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}1132 for the following reason."1133" 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-OB1.1134 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}."1135. 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.1136 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}."1137. 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 3.1138 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."1139 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 $^{+3}_{-2}$ $_{\odot}$ ."1140 We have assumed a, We have assumed a1141Effects 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)."1142 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.1143 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)."1144 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."1145 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.1146 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}.1147 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}.1148 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."1149" 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.1150" 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 $M_\star$.1151 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)."1152 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."1153" 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."1154 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}) ).1155To 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.1156 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 $R\gg R_\star$."1157 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}) )."1158" 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) $f(R)=1-(R_\star/R)^{1/2}$.1159" 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.1160" 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).1161 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.1162 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.1163 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 narrow1164spectroscopy. (e.g.Petitetal.2003)...,spectroscopy \citep[e.g.][]{2008MNRAS.388...80P}.1165 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.1166 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}."1167 We are most grateful to Dr. M. Rempel for helpful advice., We are most grateful to Dr. M. Rempel for helpful advice.1168 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.1169 The page charge for this paper is supported by CICA., The page charge for this paper is supported by CfCA.1170 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.1171"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 $w_p(r_p)$."1172" 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.1173" 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}."1174 We lind excellent agreement wilh the measurements of ?.., We find excellent agreement with the measurements of \citet{masjedi/etal:2006}.1175 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."1176" 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}) );"1177 P=29 for all 18 points., $\chi^2 = 29$ for all 18 points.1178 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.1179" 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 $w_p(r_p)$."1180" 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.1181" 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."1182" 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"1183the odf files with the SAS software (version 7.0.0).,the odf files with the SAS software (version 7.0.0).1184 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."1185 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."1186 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.1187 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."1188 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."1189 Spectra were binned according to the luminosity of each source., Spectra were binned according to the luminosity of each source.1190 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.1191 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.12.4.0.1192 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."1193 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.1194 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."1195 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).1196 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)."1197 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."1198" 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."1199 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."1200 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).

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