ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 According to our information. G89-1H with ΠΠ=—1.9 (CLLA) is the most metal-poor quadruple svstem known to date. which makes it an interesting object lor a more detailed study.," According to our information, G89-14 with $\mathrm{[m/H]}=-1.9$ (CLLA) is the most metal-poor quadruple system known to date, which makes it an interesting object for a more detailed study."3 Stellar streams (e.g.. Eggen 1996a.b)) are associations of stars possessing simular kinematies and metallicity.," Stellar streams (e.g., \citealt{eggen_1996a,eggen_1996b}) ) are associations of stars possessing similar kinematics and metallicity."4 The study of such streams allows restoring (ο a certain degree the picture of the formation of various dynamical structures in our Galaxy., The study of such streams allows restoring to a certain degree the picture of the formation of various dynamical structures in our Galaxy.5 Traditionally. the stellar streams are being selected in a certain phase space and then (heir origin is interpreted using the data of spectroscopic analvsis.," Traditionally, the stellar streams are being selected in a certain phase space and then their origin is interpreted using the data of spectroscopic analysis."6 Ii a phase space. a fine structure like stellar multiplicity can give additional information on the dynamical evolution of the stream and ils primogenitor.," In a phase space, a fine structure like stellar multiplicity can give additional information on the dynamical evolution of the stream and its primogenitor."7 However. until now il was not taken into due consideration.," However, until now it was not taken into due consideration."8 The following six stars of our sample:QG10-4..G13-9..G60-48... G24-3.. GI8-54.. G28-43.. ave part of the moving group (Eegen1996a).. 10 other objects:G130-65.. G'15-56..G5-35..G40-14.. GI14-25.. GII-4H..G13-35.. GISS-IL1.. GIS2-32.. G126-52.. belong to the moving group (Eeeen1996b).," The following six stars of our sample:, are part of the moving group \citep{eggen_1996a}, 10 other objects:, , , belong to the moving group \citep{eggen_1996b}."9. In Table G we are listing some characteristics ofthese (wo halostreams., In Table \ref{streams} we are listing some characteristics ofthese two halostreams.10 In the penultimate columnof the table. vou," In the penultimate columnof the table, you"11regime on the HR diagram.,regime on the HR diagram.12 The peak of the mass loss rate ts reached when the helium mass fraction at the centre decreases to 0.54 (0.031). and the total mass to 9.35 (6.50 Μ.Ο. for 25 (20 M.)) model sequence.," The peak of the mass loss rate is reached when the helium mass fraction at the centre decreases to 0.54 (0.031), and the total mass to 9.35 (6.50 ), for 25 (20 ) model sequence."13 As shown for the 20M sequence. such a mass loss history of runaway increase followed by sudden decrease can be repeated as the star moves in and out the unstable regime.," As shown for the 20 sequence, such a mass loss history of runaway increase followed by sudden decrease can be repeated as the star moves in and out the unstable regime."14" The 20 star finally ends its life as a yellow supergiant of Mj,26.1M... with only a small amount of hydrogen of about 0.5."," The 20 star finally ends its life as a yellow supergiant of $M_\mathrm{tot} = 6.1$, with only a small amount of hydrogen of about 0.5."15... The most likely outcome of the death of such a star would be a Type IIb supernova., The most likely outcome of the death of such a star would be a Type IIb supernova.16 The 25 star has a hydrogen envelope of 0.22 at core helium exhaustion... and would produce either a Type Ib or a IIb supernova depending on the subsequent history of mass loss.," The 25 star has a hydrogen envelope of 0.22 at core helium exhaustion, and would produce either a Type Ib or a IIb supernova depending on the subsequent history of mass loss."17 When the mass loss rate reaches the maximum. the growth rate given by Eq. (1))," When the mass loss rate reaches the maximum, the growth rate given by Eq. \ref{eq1}) )"18 becomes as high as 11.3 and 10.6 for 25 and 20 models. respectively.," becomes as high as 11.3 and 10.6 for 25 and 20 models, respectively."19 We checked tf non-linear evolutionary.. calculations also give such high growth rates. given that the relation of Eq. (1))," We checked if non-linear evolutionary calculations also give such high growth rates, given that the relation of Eq. \ref{eq1}) )"20 1s only based on the result with 1)<8.0 (Fig. 3)., is only based on the result with $\eta \la 8.0$ (Fig. \ref{fig:eta}) ).21 We find that. at the reference points marked by the filled circles in Fig. 4..," We find that, at the reference points marked by the filled circles in Fig. \ref{fig:evol},"22 the hydrodynamie calculations give jj= 9.9 and 9.0 for 25 and 20M. models respectively. which are comparable to the values given by Eq. (1)).," the hydrodynamic calculations give $\eta = $ 9.9 and 9.0 for 25 and 20 models respectively, which are comparable to the values given by Eq. \ref{eq1}) )."23 If a pulsationally-driven super-wind (PDSW) phase could be induced by strong pulsations. this would have very important implications for supernova progenitors.," If a pulsationally-driven super-wind (PDSW) phase could be induced by strong pulsations, this would have very important implications for supernova progenitors."24 As implied by our model sequences presented above. such a PDSW can significantly reduce the eritical ZAMS mass (Ma) above which a huge fraction of the stellar envelope is removed before core collapse. thus producing no Type H-P supernova.," As implied by our model sequences presented above, such a PDSW can significantly reduce the critical ZAMS mass $M_\mathrm{crit}$ ) above which a huge fraction of the stellar envelope is removed before core collapse, thus producing no Type II-P supernova."25 Based on our results. we suggest M with a PDSW (see Figs.," Based on our results, we suggest $M_\mathrm{crit} \sim$ with a PDSW (see Figs."26 | and [ 4)). while the models with the JNH88 mass loss rate predict SN II-P. progenitors up to ~....," \ref{fig:hr} and \ref{fig:evol}) ), while the models with the JNH88 mass loss rate predict SN II-P progenitors up to $\sim$."27 Interestingly. this value of M. 1s comparable to what other alternative prescriptions of RSG winds mass loss predict (Salasnichetal.1999:Vanbeveren2007).," Interestingly, this value of $M_\mathrm{crit}$ is comparable to what other alternative prescriptions of RSG winds mass loss predict \citep{Salasnich99, vanBeveren07}."28. Note that Mi could be even further reduced by rotation. as implied by the result of Hegeretal.(1997) (seealsoMeynet&Maeder 2003).," Note that $M_\mathrm{crit}$ could be even further reduced by rotation, as implied by the result of \citet{Heger97}~ \citep[see also][]{Meynet03}."29 This might provide a plausible solution to the so-called RSG problem. t.e. the observed lack of type II-P progenitors with Mii216.53:1.5M. (Smarttetal.2009).," This might provide a plausible solution to the so-called RSG problem, i.e. the observed lack of type II-P progenitors with $M_\mathrm{init} \ga 16.5 \pm 1.5$ \citep{Smartt09}."30. Such observation could also result from the presence of dusty circumstellar material around the stars. as an obscured progenitor would be estimated to have a lower initial mass in pre-SN images.," Such observation could also result from the presence of dusty circumstellar material around the stars, as an obscured progenitor would be estimated to have a lower initial mass in pre-SN images."31 However Smartt(2009) also noted that the number of SNe type II-P in their sample is consistent with the expected number of stars in the range 8.5-16.5 assuming a Salpeter IMF, However \citet{Smartt09} also noted that the number of SNe type II-P in their sample is consistent with the expected number of stars in the range 8.5-16.5 assuming a Salpeter IMF.32 This could be a possible indication of the fact that stars with higher initial mass do actually lose most of their hydrogen envelope. which cannot be easily understood with the canonical mass loss rate of JNH88.," This could be a possible indication of the fact that stars with higher initial mass do actually lose most of their hydrogen envelope, which cannot be easily understood with the canonical mass loss rate of JNH88."33 The PDSW is expected to significantly affect the circumstellar medium around a RSG (van Veelen. in prep.).," The PDSW is expected to significantly affect the circumstellar medium around a RSG (van Veelen, in prep.)."34 When the star dies. the shock produced by the collision between the SN ejecta and the circumstellar material transforms kinetic energy into thermal energy.," When the star dies, the shock produced by the collision between the SN ejecta and the circumstellar material transforms kinetic energy into thermal energy."35 This energy can be radiated away at different wavelengths. resulting in brightnening the supernova remnant for long times (Chevalier1977).," This energy can be radiated away at different wavelengths, resulting in brightnening the supernova remnant for long times \citep{Chevalier77}."36. If the collision occurs directly after the SN explosion. the emission can even alter the spectrum and light curve of the SN.," If the collision occurs directly after the SN explosion, the emission can even alter the spectrum and light curve of the SN."37 This scenario 1s usually invoked to explain the class of Type IIn supernovae (Schlegel1990:Filippenko1997).," This scenario is usually invoked to explain the sub-class of Type IIn supernovae \citep{Schlegel90,Filippenko97}."38. 'To reproduce the light curve of the most luminous Type ΠΠ supernovae (SNe IIn) like SN 2006gy and SN 2006tf. very massive shells are required. which need to be ejected in eruptive events a few years before core collapse (e.g...vanMarleetal.2010).," To reproduce the light curve of the most luminous Type IIn supernovae (SNe IIn) like SN 2006gy and SN 2006tf, very massive shells are required, which need to be ejected in eruptive events a few years before core collapse \citep[e.g., ][]{vanMarle10}."39. Pulsational pair instability (Woosleyetal.2007) and LBV-like eruptions (e.g..Smithetal.2007) have been discussed to explain the presence of shells of around the progenitors of luminous SNe IIn.," Pulsational pair instability \citep{Woosley07} and LBV-like eruptions \citep[e.g., ][]{Smith07}40 have been discussed to explain the presence of shells of around the progenitors of luminous SNe IIn."41 A PDSW phase is an unlikely explanation for such extreme environments., A PDSW phase is an unlikely explanation for such extreme environments.42 However it might be interesting for those type In where the required mass in the stellar vicinity ts of the order of a few solar masses or less., However it might be interesting for those type IIn where the required mass in the stellar vicinity is of the order of a few solar masses or less.43 In this context it is interesting to consider the circumstellar medium around the RSG VY CMa., In this context it is interesting to consider the circumstellar medium around the RSG VY CMa.44 The stellar surrounding appears shaped by episodic mass ejections which occurred about 500-1000 yr ago (Smithetal.2009).., The stellar surrounding appears shaped by episodic mass ejections which occurred about 500-1000 yr ago \citep{Smith09}. .45 The complex morphology of the CSM is suggestive of a possible interaction between convection and pulsation (seeFig.13inSmithetal.2009). which for these stars are predicted to occur on similar timescales (e.g..Hegeretal.1997).," The complex morphology of the CSM is suggestive of a possible interaction between convection and pulsation \citep[see Fig.~13 in][]{Smith09}, which for these stars are predicted to occur on similar timescales \citep[e.g., ][]{Heger97}."46. The mass loss rate of 1—2«107Myer! derived by Smithetal.(2009) is comparable to the one expected during the PDSW phase (see Fig. 4))., The mass loss rate of $1-2\times10^{-3}~\mathrm{M_\odot yr^{-1}}$ derived by \citet{Smith09} is comparable to the one expected during the PDSW phase (see Fig. \ref{fig:evol}) ).47 Even if the mass loss prescription for PDSW we used is somewhat arbitrary. we want to stress that the energy available from the growth of pulsation is enough to drive mass loss rates up to ~107M.yr!.," Even if the mass loss prescription for PDSW we used is somewhat arbitrary, we want to stress that the energy available from the growth of pulsation is enough to drive mass loss rates up to $\sim4810^{-2}~\mathrm{M_\odot yr^{-1}}$."49 Smithetal.(2009) state that an extreme RSG like VY CMa would produce a Type IIn event like SN 1988Z if it were to explode in its current state., \citet{Smith09} state that an extreme RSG like VY CMa would produce a Type IIn event like SN 1988Z if it were to explode in its current state.50 Therefore the occurrence of pulsation-driven super-winds in RSG might explain moderately luminous SN IIn if the enhanced mass loss takes place less than about ~10? years before core collapse., Therefore the occurrence of pulsation-driven super-winds in RSG might explain moderately luminous SN IIn if the enhanced mass loss takes place less than about $\sim10^3$ years before core collapse.51 With the mass loss prescription of Eq., With the mass loss prescription of Eq.52 2. with a sufficiently large à (~ 2). this would occur in a narrow range around — for non-rotating stars.," \ref{eq2} with a sufficiently large $\alpha$ $\sim 2$ ), this would occur in a narrow range around $\sim$ for non-rotating stars."53 An accurate determination of the expected.. rate of type In due to PDSW requires a more realistic mass loss prescription., An accurate determination of the expected rate of type IIn due to PDSW requires a more realistic mass loss prescription.54 To conclude. if a pulsationally-driven super-wind phase could be induced by strong pulsations. we would expect a substantial change in the late evolution of single massive stars.," To conclude, if a pulsationally-driven super-wind phase could be induced by strong pulsations, we would expect a substantial change in the late evolution of single massive stars."55 The mass loss rate during the RSG phase would increase dramatically for stars with Mini2Moi compared to the JNH88 rate. and the PDSW phaseshould start earlier in the evolution of more massive stars.," The mass loss rate during the RSG phase would increase dramatically for stars with $M_\mathrm{init} \ga M_\mathrm{crit}$ compared to the JNH88 rate, and the PDSW phaseshould start earlier in the evolution of more massive stars."56 The resulting pre-supernova structure of these stars is affected. as well as their CSM.," The resulting pre-supernova structure of these stars is affected, as well as their CSM."57 For single stars. this suggests the following sequence in supernova types as function of increasing initial mass: P>I-L ," For single stars, this suggests the following sequence in supernova types as function of increasing initial mass: II-P $\xrightarrow{}$ "58N-rayv source (Isracl&Stella1996:Vaughan2005).,"X-ray source \citep{Israel96,Vaughan05}."59. We fitted the power versus frequency relation for periods im the range 6280 davs to a power-law form aud found a spectral index of —1.01dk0.06., We fitted the power versus frequency relation for periods in the range 6–280 days to a power-law form and found a spectral index of $-1.04 \pm 0.06$.60 We generated red noise with a spectral index of —1.10 and with mean aud variance equal to those of the data using the routine of the IDL subroutine library provided by the Institut fiürr Astronomie und Astroplisik of the Universitàtt Tibbinecu (Timmer&Konig1991)., We generated red noise with a spectral index of $-1.10$ and with mean and variance equal to those of the data using the routine of the IDL subroutine library provided by the Institut fürr Astronomie und Astrophysik of the Universitätt Tübbingen \citep{Timmer94}.61 The duration of each generated light curve is longer than the actual data im order to minimize the effects of red noise leakage., The duration of each generated light curve is longer than the actual data in order to minimize the effects of red noise leakage.62 Each light curve coutains 8192 data points with uniform spacing of 0.66 days and a subset of 330 points from the muddle of this set. with relative times matching the actual observations are extracted for analysis., Each light curve contains 8192 data points with uniform spacing of 0.66 days and a subset of 330 points from the middle of this set with relative times matching the actual observations are extracted for analysis.63 These 330 simulated fluxes were processed with the same procedures used to analyze the real cata., These 330 simulated fluxes were processed with the same procedures used to analyze the real data.64 We generated 2«10° trial Πο curves aud searched for cases where the power at periods of 10 to 150 days was greater than or equal to the observed value of 77.9., We generated $2 \times 10^{6}$ trial light curves and searched for cases where the power at periods of 10 to 150 days was greater than or equal to the observed value of 77.9.65 We found one such case and estimate the probability of chance ocemrence of our obscrved signal to be 5«10., We found one such case and estimate the probability of chance occurrence of our observed signal to be $5 \times 10^{-7}$.66* Fitting the distribution of maxima observed powers. we estimate that the probability of chance occurrence of our observed signal is 6«LO*. iu good agreement.," Fitting the distribution of maximum observed powers, we estimate that the probability of chance occurrence of our observed signal is $6 \times 10^{-7}$, in good agreement."67 This procedure is conservative because it includes the signal in the caleulation of the power-law slope aud he variance aud because the period search range. 10150 davs. extends to significantly lower frequencies than he observed period where the red noise produces high amplitude fluctuations.," This procedure is conservative because it includes the signal in the calculation of the power-law slope and the variance and because the period search range, 10--150 days, extends to significantly lower frequencies than the observed period where the red noise produces high amplitude fluctuations."68 If we restrict the search range to veriods of 62 davs or less. then a fit to the distribution of the maxiuun observed powers indicates that the xobabilitv of chance occurrence of our observed. signal i310P.," If we restrict the search range to periods of 62 days or less, then a fit to the distribution of the maximum observed powers indicates that the probability of chance occurrence of our observed signal is $3 \times 10^{-13}$."69 The coliereuce or quality value of the peak signal. the oeriod of the peak divided by the full width at half uaxinmun power. is Q=22.3.," The coherence or quality value of the peak signal, the period of the peak divided by the full width at half maximum power, is $Q = 22.3$."70 This is fully cousisteut with that expected for a periodic process eiven the observation curation., This is fully consistent with that expected for a periodic process given the observation duration.71 Fig., Fig.72 3. shows the data folded at the best fit period., \ref{phase} shows the data folded at the best fit period.73 Each point is the average flux of the observations falling within the eiven phase bin aud the error bar is the standard error. the standard deviation of the fluxes in each bin divided by the square root of the ΙΟ of fluxes im that bin.," Each point is the average flux of the observations falling within the given phase bin and the error bar is the standard error, the standard deviation of the fluxes in each bin divided by the square root of the number of fluxes in that bin."74" The amplitude of the modulation. taken as the miaxinmni average flux in one bin minus the πα, is (0.99£0.10)ς1οHerecni24ot."," The amplitude of the modulation, taken as the maximum average flux in one bin minus the minimum, is $(0.99 \pm750.10) \times 10^{-11} \rm \, erg \, cm^{-2} \, s^{-1}$."76 To search for rapid variability. we extracted event files with high time resolution data for the 187 observations.," To search for rapid variability, we extracted event files with high time resolution data for the 187 observations."77 Eveuts in the 2.1-11.9 keV energy. band were selected in the good time intervals defined above aud split iuto segments of 256 x. Au FFT with a time resolution of ls was calculated for cach scemenut., Events in the 2.4-11.9 keV energy band were selected in the good time intervals defined above and split into segments of 256 s. An FFT with a time resolution of 1 s was calculated for each segment.78 The FFTs within cach observation were added iucolercutly., The FFTs within each observation were added incoherently.79 The resulting, The resulting80"As described in the introduction, observational evidence for vortical motion on the solar surface typically refers to larger scales than the strong, small-scale vortices studied here, although the general characteristics are similar (e.g., the association with downflows).","As described in the introduction, observational evidence for vortical motion on the solar surface typically refers to larger scales than the strong, small-scale vortices studied here, although the general characteristics are similar (e.g., the association with downflows)."81 The reported mean lifetimes of granular-scale vortices of 5-8 minutes (??) are not drastically different from the value of 3.5min estimated above.," The reported mean lifetimes of granular-scale vortices of 5–8 minutes \citep{2008Bonet,2010Bonet} are not drastically different from the value of $3.5\unit{min}$ estimated above."82" However, because their rotation periods are different, our vortices make about 2 revolutions during their lifetimes, while the observed vortices can be followed for only a fraction of one rotation (forinstance,about2596ofaperiodinthecase?).."," However, because their rotation periods are different, our vortices make about 2 revolutions during their lifetimes, while the observed vortices can be followed for only a fraction of one rotation \citep[for instance, about 25\% of83a period in the case of][]{2010Bonet}."84 It is conceivable that the observed vortical motions represent the peripheral parts of the much stronger small-scale vortex cores that show up in the simulations but are too small to be observed directly., It is conceivable that the observed vortical motions represent the peripheral parts of the much stronger small-scale vortex cores that show up in the simulations but are too small to be observed directly.85 The outer vortex parts would be much more strongly affected by the evolving granulation pattern and thus be detectable only for a fraction of a rotation period., The outer vortex parts would be much more strongly affected by the evolving granulation pattern and thus be detectable only for a fraction of a rotation period.86" To see whether the vortices studied in this paper would be detectable in observational data through feature tracking techniques, we consider horizontal pseudo pathlines (trajectories of fluid elements) in Fig. 13.."," To see whether the vortices studied in this paper would be detectable in observational data through feature tracking techniques, we consider horizontal pseudo pathlines (trajectories of fluid elements) in Fig. \ref{fig:cork}."87 The pathlines are determined from 10 snapshots of the horizontal velocity field at the average height of the optical surface with a temporal spacing of )., The pathlines are determined from 10 snapshots of the horizontal velocity field at the average height of the optical surface with a temporal spacing of.88".T heyare'pseudo"" becausetheverticalvelocityisignored.", They are “pseudo” because the vertical velocity is ignored.89" Some pathlines are spiraling in towards a region of high swirling strength, see the feature at )intheleft -handpanelo fFig. 13."," Some pathlines are spiraling in towards a region of high swirling strength, see the feature at in the left-hand panel of Fig. \ref{fig:cork}."90".Inthiscase, thereisaclearassociationwit scalevorticespresentedinthispaper."," In this case, there is a clear association with the small-scale vortices presented in this paper."91"However, wealso findpathlineswhich upregionshowninFig. 13.."," However, we also find pathlines which are curved but for which the association with an actual vortex is not clear, for instance in the blown-up region shown in Fig. \ref{fig:cork}."92 We have investigated vortical fluid motions in simulations of near-surface solar convection by calculating the eigenvalues and eigenvectors of the velocity gradient tensor field., We have investigated vortical fluid motions in simulations of near-surface solar convection by calculating the eigenvalues and eigenvectors of the velocity gradient tensor field.93 Complex eigenvalues with a large imaginary part indicate regions of strong swirling., Complex eigenvalues with a large imaginary part indicate regions of strong swirling.94" They are found predominantly in and near the intergranular lanes, where cooled fluid is sinking down in a turbulent fashion."," They are found predominantly in and near the intergranular lanes, where cooled fluid is sinking down in a turbulent fashion."95" The swirling regions form an unsteady network of highly tangled filaments, some of which protrude above the optical surface."," The swirling regions form an unsteady network of highly tangled filaments, some of which protrude above the optical surface."96" Near the optical surface, vertically oriented swirls are preferentially located in the interior of intergranular lanes, where the downflow is strong."," Near the optical surface, vertically oriented swirls are preferentially located in the interior of intergranular lanes, where the downflow is strong."97" Horizontal swirls, on the other hand, are predominantly located at the edges of the granules, where vertical motion is mostly absent."," Horizontal swirls, on the other hand, are predominantly located at the edges of the granules, where vertical motion is mostly absent."98" The 3D structure of contiguous features above the optical surface is manifold, but often in the form of bent and arc-shaped filaments."," The 3D structure of contiguous features above the optical surface is manifold, but often in the form of bent and arc-shaped filaments."99" These type of structures have previously been seen in independent numerical simulations with different codes, notably ? and ?.."," These type of structures have previously been seen in independent numerical simulations with different codes, notably \citet{2010Muthsam} and \citet{1998Stein}."100 The swirling direction (rotation axis) is typically aligned with the longitudinal direction(s) in a contiguous feature., The swirling direction (rotation axis) is typically aligned with the longitudinal direction(s) in a contiguous feature.101width observed here. show that the effect described above is unimportant for IHT2O masers.,"width observed here, show that the effect described above is unimportant for $_2$ O masers."102 The line widths show that in almost all cases gQὃνFis satistied., The line widths show that in almost all cases $g\Omega \gg R$ is satisfied.103 Wiebe Watson (1998) have shown that the wropagation of linear polarization can also resul circular polarization., Wiebe Watson (1998) have shown that the propagation of linear polarization can also result in circular polarization.104 For linear polarizatious of z105 he resulting circular polarization is of the same order of naguitude as the polarization due to the regular Zeca interpretation. while maeuetic fields could again be a actor of 1000 less.," For linear polarizations of $\approx 10\%$, the resulting circular polarization is of the same order of magnitude as the polarization due to the regular Zeeman interpretation, while magnetic fields could again be a factor of 1000 less."105 Below we indicate that also this 1011”Zeenau interpretation is unlikely duc to the lack of linear volarization., Below we indicate that also this non-Zeeman interpretation is unlikely due to the lack of linear polarization.106 Iu the case of the LTE οσα models we first fit a conibination of the 3 strongest lywperfine components to the total power spectrin., In the case of the LTE Zeeman models we first fit a combination of the 3 strongest hyperfine components to the total power spectrum.107 From this we ect the line width Acy of the maser feature aud the contribution to the line of the πηραπο compoucuts., From this we get the line width $\Delta v_{\rm L}$ of the maser feature and the contribution to the line of the hyperfine components.108 For this combination we calculate a svuthetic V-spectruui. such as shown in Fie.," For this combination we calculate a synthetic V-spectrum, such as shown in Fig."109 1 for the separate hvperfue compoucuts., \ref{vs} for the separate hyperfine components.110 The svuthetic spectrum is then fitted to the circular polarization spectrum., The synthetic spectrum is then fitted to the circular polarization spectrum.111 As our observations in VOL have shown that the observed. V-spectruui can be narrower than the svuthetic Vespectrmm. we also allow for the narrowing of our svuthetic spectra in the ft.," As our observations in V01 have shown that the observed V-spectrum can be narrower than the synthetic V-spectrum, we also allow for the narrowing of our synthetic spectra in the fit."112 Furthermore we need to remove the scaled down replicas of the total intensity., Furthermore we need to remove the scaled down replicas of the total intensity.113 We fit the following fiction: Here Av is the velocity in [kiu/s} measured from the peak of the total intensity. V is the svuthetic spectra Vieurk.," We fit the following function: Here $\Delta v$ is the velocity in [km/s] measured from the peak of the total intensity, $V^*$ is the synthetic spectrum $V_{\rm114synth}$ ."115 We fi for the parameters ay. e» and a3.," We fit for the parameters $a_1$, $a_2$ and $a_3$."116 The paracter ay controls the amplitude of he V-spectruni. while e» determines the amount of narrowing of the observed spectra with respect to the model spectra.," The parameter $a_1$ controls the amplitude of the V-spectrum, while $a_2$ determines the amount of narrowing of the observed spectra with respect to the model spectra."117" Finally, the removal of the scaled down total iuteusity profile is determined by a3."," Finally, the removal of the scaled down total intensity profile is determined by $a_3$."118" The auplituce e4 of the best fitted model determines Vi, aud Vias."," The amplitude $a_1$ of the best fitted model determines $V_{\rm min}$ and $V_{\rm119max}$."120 Together with the peak intensity of the maser feature lis gives Py., Together with the peak intensity of the maser feature this gives $P_{\rm V}$.121 With the previously determined line width ο we calculate By using Eq. L..," With the previously determined line width $\Delta122v_{\rm L}$ we calculate $B_{||}$ using Eq. \ref{eq2}."123 For the nou-LTE case. we fit out inodels to both total intensity aud circular polarization simultaneously.," For the non-LTE case, we fit out models to both total intensity and circular polarization simultaneously."124 We use Eq., We use Eq.125 3 with VVinsda," \ref{eq4} with $V^* =126V_{\rm model}$."127 Because the noi-LTE models are already intrinsically iuore uarrow. we first fik the width of the V-spectruu by setting ου=1.0 and only fitting ο aud a3.," Because the non-LTE models are already intrinsically more narrow, we first fix the width of the V-spectrum by setting $a_2128= 1.0$ and only fitting $a_1$ and $a_3$."129" We have chosen to fit models with Crheaual0.58 and 1.0 kms. as the line widths of the maser features indicate that wach higher iutziusic thermal widths are unlikely,"," We have chosen to fit models with $v_{\rm thermal} = 0.8$ and $1.0$ km/s, as the line widths of the maser features indicate that much higher intrinsic thermal widths are unlikely."130 We fiud that iu some cases the circular polarization spectrum is still narrower than the model spectrin., We find that in some cases the circular polarization spectrum is still narrower than the model spectrum.131 For these features we allow for narrowing bx releasing e»., For these features we allow for narrowing by releasing $a_2$.132 The combined fitting determines the circular polarization percentage Py aud the line width Acer., The combined fitting determines the circular polarization percentage $P_{\rm V}$ and the line width $\Delta v_{\rm L}$.133 The coefficient “lpp is specifically determined for the best fitted model., The coefficient $A_{\rm F-F'}$ is specifically determined for the best fitted model.134 Using this Eq., Using this Eq.135 L again gives By., \ref{eq2} again gives $B_{||}$.136 The observations were performed at the Very Loug baseline Array (VLBA) on December 13th 1998., The observations were performed at the Very Long baseline Array (VLBA) on December 13th 1998.137" At the frequency of the 6,5555 rotational transition of ITO. 22.235 (11. the average beam width is 2:0.5< mas."," At the frequency of the $6_{16}-5_{23}$ rotational transition of $_{2}$ O, 22.235 GHz, the average beam width is $\approx 0.5 \times1380.5$ mas."139 This allows us to resolve the different TeO maser features in the CSE., This allows us to resolve the different $_{2}$ O maser features in the CSE.140 The data were correlated twice. ounce with modest (7.8 kIIz 20.1 lans. +) spectral resolution. which enabled us to generate all Ll polarization combinations (RR. LL. RE aud LR).," The data were correlated twice, once with modest $7.8$ kHz $= 0.1$ km $^{-1}$ ) spectral resolution, which enabled us to generate all 4 polarization combinations (RR, LL, RL and LR)."141 The secoud correlator run was performed with high spectral resolution (1.95 kIIz =0.027 kii C) necessary fo detect the circular polarization signature of the TeO Zeeman splitting. aud therefore ouly contained the two polarization combinations RR aud LL.," The second correlator run was performed with high spectral resolution $1.95$ kHz $= 0.027$ km $^{-1}$ ), necessary to detect the circular polarization signature of the $_{2}$ O Zeeman splitting, and therefore only contained the two polarization combinations RR and LL."142 We have performed 6 hours of observations per source-calibrator pair., We have performed 6 hours of observations per source-calibrator pair.143 The calibrator was observed for 11/2 hours imm a number of scans equally distributed over the 6 hours., The calibrator was observed for $1~1/2$ hours in a number of scans equally distributed over the 6 hours.144 We used 2 filters (TFs) of 1 MIIz width. which were overlapped to get a velocity coverage of z22 laus. This covers most of the velocity range of the Πο Ο maser.," We used 2 filters (IFs) of 1 MHz width, which were overlapped to get a velocity coverage of $\approx 22$ km/s. This covers most of the velocity range of the $_2$ O maser."145 The data analvsis followed the method. of Keimball. Diamond Cotton (1995).," The data analysis followed the method of Kemball, Diamond Cotton (1995)."146" The reduction path. shown in Fig. δν,"," The reduction path, shown in Fig. \ref{fig1},"147 was performed in the AIPS data reduction package., was performed in the AIPS data reduction package.148 The first standard calibration steps were performed on the data-set with modest spectral resolution., The first standard calibration steps were performed on the data-set with modest spectral resolution.149 We used the svsteui temperature measurements provided with the data to perform the amplituce calibration for both the calibrators aud the sources., We used the system temperature measurements provided with the data to perform the amplitude calibration for both the calibrators and the sources.150 Parallactic auele correction. flageineao aud sinele- aud imultibanud delay calibration were all oue τοσα] on the calibrators observed with each source.," Parallactic angle correction, flagging and single- and multi-band delay calibration were all done regularly on the calibrators observed with each source."151" During 1/3 of the observatiou tie, he first IF suffered strong interference. which forced us to Πας several of the frequency clamuels (àL1 MITZ) iu addition to the normal flaeeime."," During 1/3 of the observation time, the first IF suffered strong interference, which forced us to flag several of the frequency channels $\approx 0.1$ MHz) in addition to the normal flagging."152 Also. for most of the observations IF 1 of the Los Alamos (LÀ) auteuna was nutsable.," Also, for most of the observations IF 1 of the Los Alamos (LA) antenna was unusable."153 The solutions obtained at these calibration steps were copied aud applied to the high spectral resolution data set., The solutions obtained at these calibration steps were copied and applied to the high spectral resolution data set.154 The complex bandpasses were then determined for bothdata-sets separately., The complex bandpasses were then determined for bothdata-sets separately.155 Additional calibration steps were needed for accurate processing of polarization data., Additional calibration steps were needed for accurate processing of polarization data.156 The gain ratio between the R- (vight-}) and L- hand polarizations was determined using the auto- data of the reference antenna on a short scan of the maser source., The gain ratio between the R- (right-) and L- (left-) hand polarizations was determined using the auto-correlation data of the reference antenna on a short scan of the maser source.157 This step contains the critical, This step contains the critical158thus. the line of sight is unbiased for intervening-metal line absorbers.,"thus, the line of sight is unbiased for intervening–metal line absorbers."159 The stronger absorption proliles are resolved. auc required. a two-component Gaussian fit., The stronger absorption profiles are resolved and required a two–component Gaussian fit.160 A2803 15 blended. with Galactic A2852: this resulted in an unphysical doublet ratio for the red component of the Gaussian fit., $\lambda 2803$ is blended with Galactic $\lambda 2852$; this resulted in an unphysical doublet ratio for the red component of the Gaussian fit.161 Ouly a C270H spectrum was available for this line of sight., Only a G270H spectrum was available for this line of sight.162 Because of the relatively high redshift of this emission liue object. the strougFell trausitious fell in tle [orest.," Because of the relatively high redshift of this emission line object, the strong transitions fell in the forest."163 A2600 is Just recdhward of the emission. but it is hopelessly bleucded in a strong absorption complex.," $\lambda 2600$ is just redward of the emission, but it is hopelessly blended in a strong absorption complex."164 À2822 is quite strong.," $\lambda1652852$ is quite strong."166 The data are preseuted iu Figure 3. and the rest-frame equivalent wicltlis are presented in Table 2., The data are presented in Figure \ref{fig:q1327} and the rest–frame equivalent widths are presented in Table 2.167 This object was observed for a program (PID 6781) to investigate the iouizine coutinuum in ACUN: this is an unbiased line of sight., This object was observed for a program (PID 6781) to investigate the ionizing continuum in AGN; this is an unbiased line of sight.168 The system is the weakest of the four ancl the ouly one in which the profiles are uuresolved at FOS resolution (subsequently. the widths of the Craussians used for the equivalent width measurements were held coustant at the value of the FOS instrumental spreacl Πιοτοι: this leaves a somewhat significaut residual to the fit iu the liue cores).," The system is the weakest of the four and the only one in which the profiles are unresolved at FOS resolution (subsequently, the widths of the Gaussians used for the equivalent width measurements were held constant at the value of the FOS instrumental spread function; this leaves a somewhat significant residual to the fit in the line cores)."169 Both a GI30H and a CU90H spectrum were available for this line of sight., Both a G130H and a G190H spectrum were available for this line of sight.170 The rest-frame equivalent width is 3.05x0.26A. indicatiug that this systems is not a DLA.," The rest–frame equivalent width is $3.95\pm0.26$, indicating that this systems is not a DLA."171 No nor absorption was detected to a 5o detection limit of 0.27 and 0.36. A. respectively.," No nor absorption was detected to a $5~\sigma$ detection limit of $0.27$ and $0.36$ , respectively."172 AI331 was detected., $\lambda 1334$ was detected.173 Also detected were A2301. 2371. 2382. ancl 2600.," Also detected were $\lambda 2344$, 2374, 2382, and 2600."174 A2852 was not detected to 0.19 (5 σ]., $\lambda 2852$ was not detected to 0.49 $5~\sigma$ ).175 The data are presented in Figure | aud the rest-frame equivalent widths are presented in Table 2., The data are presented in Figure \ref{fig:pg1427} and the rest–frame equivalent widths are presented in Table 2.176 Iu the spectrum of PISS O151039 (Boisséetal.1998).. absorption at 2=0.072 is associated with a post-star burst dwarf galaxy. (Steidel.Dickinson.&Bowen1993).," In the spectrum of PKS $0454+039$ \citep{boisse}, absorption at $z=0.072$ is associated with a post–star burst dwarf galaxy \citep{ccs0454}."177. Bolssé state that the presence of this galaxy was “one additional motivation for observing this quasar”., Boissé state that the presence of this galaxy was “one additional motivation for observing this quasar”.178 Thus. this quasar was dropped Crom this survey.," Thus, this quasar was dropped from this survey."179 Bowen.Blades.&Pettini(1996) published this system. which arises in the inclined galaxy M61.," \citet{bbp96} published this system, which arises in the inclined galaxy M61."180" This quasar was observed because it was a sightliue that ""lies fortuitously behind [agalaxy] whose existenceisalready known.""", This quasar was observed because it was a sightline that “lies fortuitously behind [agalaxy] whose existenceisalready known.”181 Thus. thisline of sielt was also dropped from this survey.," Thus, thisline of sight was also dropped from this survey."182The clustering of sources broadens the flux PDF and creates a value tail in the distribution (e.g. ??)).,"The clustering of sources broadens the flux PDF and creates a high-value tail in the distribution (e.g. \citealt{MD08, Dijkstra08}) )."183 How strong is this effect and how much do the properties of the underlying sources affect these trends?, How strong is this effect and how much do the properties of the underlying sources affect these trends?184" In Fig. 4,,"," In Fig. \ref{fig:cluster},"185" we plot flux PDFs at z=10, assuming Aurp Mpc."," we plot flux PDFs at $z=10$, assuming $\lmfp=10$ Mpc."186" The dotted curve in the figure was generated from the semi-numeric calculation assuming a step-function attenuation at Amtp=10 Mpc, shown to be an adequate approximation in Fig. 2.."," The dotted curve in the figure was generated from the semi-numeric calculation assuming a step-function attenuation at $\lmfp=10$ Mpc, shown to be an adequate approximation in Fig. \ref{fig:mfp_att}."187" The dashed curve is obtained in the same manner, but with randomized source locations."," The dashed curve is obtained in the same manner, but with randomized source locations."188 We see that ignoring source clustering can severely underestimate the widths of the flux PDF., We see that ignoring source clustering can severely underestimate the widths of the flux PDF.189 Even the shape of the curves is quite different., Even the shape of the curves is quite different.190 Fig., Fig.191 4 also compares the PDF to our analytic model., \ref{fig:cluster} also compares the PDF to our analytic model.192" The dot-dashed line uses the analog of equation (2)) but for step-wise attenuation (see ? and ?)), again assuming Poisson distributed sources."," The dot-dashed line uses the analog of equation \ref{eq:jdist_r0}) ) but for step-wise attenuation (see \citealt{Zuo92} and \citealt{Furlanetto09}) ), again assuming Poisson distributed sources."193 It closely matches the semi-numeric results for the same assumptions., It closely matches the semi-numeric results for the same assumptions.194 The solid line includes linear clustering over the scale Amsp as described in equation (4))., The solid line includes linear clustering over the scale $\lambda_{\rm mfp}$ as described in equation \ref{eq:jdistbn_cluster}) ).195" The clustering length of 10?Mg halos is ~ 2 Mpc at this redshift, so the clustering on scales larger than the m.f.p."," The clustering length of $M \sim \Mmin \sim 10^8 \Msun$ halos is $\sim$ 2 Mpc at this redshift, so the clustering on scales larger than the m.f.p."196 should be well into the linear regime and accurately predicted by the analytic model., should be well into the linear regime and accurately predicted by the analytic model.197" Obviously, this simple prescription provides a relatively poor match to the full semi-numeric results: although the large-scale clustering does broaden the distribution by a comparable amount, the shapes still disagree."," Obviously, this simple prescription provides a relatively poor match to the full semi-numeric results: although the large-scale clustering does broaden the distribution by a comparable amount, the shapes still disagree."198 Part of this difference is easy to understand and is a consequence of the analytic model's simple prescription for clustering: the model not only fixes the attenuation volume, Part of this difference is easy to understand and is a consequence of the analytic model's simple prescription for clustering: the model not only fixes the attenuation volume199 , 200(intermediate mass) black holes in the Galactic center region.,(intermediate mass) black holes in the Galactic center region.201 Our current observations support limits between those plotted bv Hansen&Milosavljevió(2003) in their Fig., Our current observations support limits between those plotted by \citet{HM03} in their Fig.202" 2 for “astrometric resolutions"" of 0.1 and 1.0 mas.", 2 for “astrometric resolutions” of 0.1 and 1.0 mas.203" The limits on (hese parameters are complex and depend on both parameters. but roughly we can exclude secondary. black holes with niasses greater than ~LO! and senmi-major axes between LO? and 10"" AU [romÀ*."," The limits on these parameters are complex and depend on both parameters, but roughly we can exclude secondary black holes with masses greater than $\sim10^4$ and semi-major axes between $10^3$ and $10^5$ AU from."204. Excluding stellar mass black holes (<100 )) will require more than an order of magnitude better astrometric accuracy and is unlikely in the near future., Excluding stellar mass black holes $<100$ ) will require more than an order of magnitude better astrometric accuracy and is unlikely in the near future.205 We have measured the position of the compact non-thermal radio source.A... al the center of the Galaxy relative to extragalactic radio sources.," We have measured the position of the compact non-thermal radio source, at the center of the Galaxy relative to extragalactic radio sources."206 The apparent motion of is consistent with that expected. [rom the orbit of the Sun around the Galactic center., The apparent motion of is consistent with that expected from the orbit of the Sun around the Galactic center.207 Ànv peculiar motion of perpendicular to the plane of the Galaxy is less than L8 (20)., Any peculiar motion of perpendicular to the plane of the Galaxy is less than 1.8 $2\sigma$ ).208 This result is complementary to infrared observations of stellar orbits at the Galactic center. which require 4x10° within a radius of 100 AU ofA*.," This result is complementary to infrared observations of stellar orbits at the Galactic center, which require $4\times10^6$ within a radius of 100 AU of."209. The results of several different analvses indicate a significant fraction. if not all. of the mass in ihe central LOO AU is tied toA*.," The results of several different analyses indicate a significant fraction, if not all, of the mass in the central 100 AU is tied to."210. Were not a SABIL. it must be bound to the inner 4 AU of the dvnamical center of the Galaxy.," Were not a SMBH, it must be bound to the inner 4 AU of the dynamical center of the Galaxy."211 This would imply an extraordinarily high mass density. ancl probably require a SMDII., This would imply an extraordinarily high mass density and probably require a SMBH.212 The gravitational attractions of the ~10to10* stars within 2 pe of the Galactic center impart a significant motion toÀ*., The gravitational attractions of the $\sim10^6~{\rm to}~10^7$ stars within 2 pc of the Galactic center impart a significant motion to.213. Based on nunerical simulations of the central star cluster and our upper limit to the motion of out of the plane of the Galaxy. a maxinnun-likelihood lower limit [or the mass of is 0.4x10AL.," Based on numerical simulations of the central star cluster and our upper limit to the motion of out of the plane of the Galaxy, a maximum-likelihood lower limit for the mass of is $0.4\times10^6$."214.. These analyses make very conservative assumptions that would tend to underestimate the motion of and. hence. underestimate (he mass limit.," These analyses make very conservative assumptions that would tend to underestimate the motion of and, hence, underestimate the mass limit."215 This is the first evidence (hat a compact radiative source al (he center of a galaxy is a super-lassive object., This is the first evidence that a compact radiative source at the center of a galaxy is a super-massive object.216 Other measurements determine a large mass. but can onlyindirectly associate il wilh the radiative source through positional agreement.," Other measurements determine a large mass, but can only associate it with the radiative source through positional agreement."217 The observed radio frequency size of is less than 1 AU. after accounting for the elects of interstellar scattering.," The observed radio frequency size of is less than 1 AU, after accounting for the effects of interstellar scattering."218 The mass density implied by having at least 0.4x109 within a 0.5 AU radius is a staggering 7x107 7!, The mass density implied by having at least $0.4\times10^6$ within a 0.5 AU radius is a staggering $7\times10^{21}$ $^{-3}$!219 This is only about 3 lower than the mass density of a 4x109 black hole within its Schwarzschild, This is only about 3 orders-of-magnitude lower than the mass density of a $4\times10^6$ black hole within its Schwarzschild220"profile reseiubling a star rotating at οsni=Uskinss ! (ervey dashed line, which is identical to the widest profile shown in veffiie:Sun)).","profile resembling a star rotating at $v\,\sin{i} = 4$ $^{-1}$ (grey dashed line, which is identical to the widest profile shown in \\ref{fig:Sun}) )."221 The profile expected frou: a star rotating atf οκ;= dkknnss to is a very eood match to the profile of Sanders 1152. with only the far wings bevond +6kkuiss ! differing slightly.," The profile expected from a star rotating at $v\,\sin{i} = 4$ $^{-1}$ is a very good match to the profile of Sanders 1452, with only the far wings beyond $\pm 6$ $^{-1}$ differing slightly."222 This again may be an effect of the image slicer. but the far wines of the decouvolved broadening fiction are also less well defined due to the existence of uncaptured nue bleuds and limited SNR in the data.," This again may be an effect of the image slicer, but the far wings of the deconvolved broadening function are also less well defined due to the existence of uncaptured line blends and limited SNR in the data."223 From the close match between the broadening function of Sanders 1152 and the artificially broadened versions of the Camvimede spectruii. we derive a value for the projected rotation velocity of Sanders 1152 of esiu;={4£0.5kkinss +.," From the close match between the broadening function of Sanders 1452 and the artificially broadened versions of the Ganymede spectrum, we derive a value for the projected rotation velocity of Sanders 1452 of $v\,\sin{i} = 4 \pm 0.5$ $^{-1}$."224 The uncertaiuty is an eiipirical estimate deduced. frou a comparison with broadcuing functions with different values of esin;.," The uncertainty is an empirical estimate deduced from a comparison with broadening functions with different values of $v\,\sin{i}$."225 Broadening fuuctious witli esiné=3.5 and L5kkniss+ can be distinguished from the data.," Broadening functions with $v\,\sin{i} = 3.5$ and $^{-1}$ can be distinguished from the data."226 The results of our detailed spectroscopic analysis enable us to discuss the rauge of chromospheric activity seen in M67 in conjunction with the additional -information on rotation. at least for à representative subset of solu-tvpoe stars in M67.," The results of our detailed spectroscopic analysis enable us to discuss the range of chromospheric activity seen in M67 in conjunction with the additional information on rotation, at least for a representative subset of solar-type stars in M67."227 We therefore briefly iscuss du the following some relevant issues in this contest. namely. (1) the potential rauge of brightuess variability in sun-like stars at solar age and. by inference. in the Sun itself. and (2) some facets of angular monentun evolution m solu-tvpo stars. including the possible origin of relatively rapid rotation at the age of the M67 open cluster.," We therefore briefly discuss in the following some relevant issues in this context, namely, (1) the potential range of brightness variability in sun-like stars at solar age and, by inference, in the Sun itself, and (2) some facets of angular momentum evolution in solar-type stars, including the possible origin of relatively rapid rotation at the age of the M67 open cluster."228 The potential excursions of the activity. cvcles of the M67 solu-like stars and possibly the Sun itself to exceptionally high values. as inferred from the II distvibution of the M67. solu-tvpo stars given iu Ciunupapaetal(2006.theirFie.Όλι now must be considered in the light of the rotation measures and estimates eiven herein.," The potential excursions of the activity cycles of the M67 solar-like stars and possibly the Sun itself to exceptionally high values, as inferred from the HK distribution of the M67 solar-type stars given in \citet[][their Fig. 3]{Giampapa06}, now must be considered in the light of the rotation measures and estimates given herein."229 In particubu. since 81152 with a meu HE iudex of Ll is rotating at more than twice the equatorial solar rotation velocity. aud $8717 with a mean HER = 351mA is a spectroscopic binary. the implication is that excursions iu the eveles of M67 solar-type stars aud the Suu itself appear to be less than about TIS ~ 250mA. re. the next highest ITI iudex found iu the Cdammpapa et al.," In particular, since S1452 with a mean HK index of 414 is rotating at more than twice the equatorial solar rotation velocity, and S747 with a mean HK = 354 is a spectroscopic binary, the implication is that excursions in the cycles of M67 solar-type stars and the Sun itself appear to be less than about HK $\sim$ 250, i.e., the next highest HK index found in the Giampapa et al."230 sample. which is roughly higher than the representative riaxiniuun value seen in the modern solar evele of HIS. z 225mA.," sample, which is roughly higher than the representative maximum value seen in the modern solar cycle of HK $\approx$ 225."231 Within the reported BW color range of the Sum of about 0.63 to 0.67 (VandeuBere&Dridees198£).. the only star iu the Cdaimpapa et al.," Within the reported $B-V$ color range of the Sun of about 0.63 to 0.67 \citep{VandenBerg84}, the only star in the Giampapa et al."232 sample that exceeds the maxima solar III iudex (8 1011) also is a short-period biuuv with a period of 16.2 davs aud mean IIR. = 218mA., sample that exceeds the maximum solar HK index (S 1014) also is a short-period binary with a period of 16.2 days and mean HK = 248.233 The determination of definitive τηραΠΠ to the ITI iudex for single sun-like astars at solar age aud metallicitv will require a more extensive survev of rotation in the MIG? sobu-tvpo stars., The determinination of a definitive upper-limit to the HK index for single sun-like stars at solar age and metallicity will require a more extensive survey of rotation in the M67 solar-type stars.234 Caven our results and the well-known correlation between variations in chromospheric ciission aud changes in the solar radiance or. correspoucinely. ii stellar brightness. it is of interest to consider the implications of the plausible upper liit to chromospheric emissiou in solar-type aud solu-age stars. as inferred from the M67. sample. for the possible rauge of brightuess variability that nav occur.," Given our results and the well-known correlation between variations in chromospheric emission and changes in the solar irradiance or, correspondingly, in stellar brightness, it is of interest to consider the implications of the plausible upper limit to chromospheric emission in solar-type and solar-age stars, as inferred from the M67 sample, for the possible range of brightness variability that may occur."235 Using the dependent calibration approach described i Cdampapa and adopting a color of B.V=0.65 as representativo of analogs of the Sun. we find that the upper Init eiven above of IT& z 250 corresponds to log Rope~LO. where Πέις is the ratio of the total chromospheric IT aud Is emission core flux to the stellar bolometric flux. corrected for the non-chromospheric (plotospheric) contribution.," Using the color-dependent calibration approach described in \citet{Giampapa06} and adopting a color of $B-V = 0.65$ as representative of analogs of the Sun, we find that the upper limit given above of HK $\simeq$ 250 corresponds to log $^{\prime}_{HK} \sim -4.70$, where $^{\prime}_{HK}$ is the ratio of the total chromospheric H and K emission core flux to the stellar bolometric flux, corrected for the non-chromospheric (photospheric) contribution."236 Inspection of the results of long-term. high precision photometry of solu-tvpe stars by Lockwoodetal.(19907.theirFig.17) as a muction of the paramcter Πτι suggests that this level of activity would correspond to an annual mean level of rum brightuess variations of roughly 0.002 mae. ie. varlabilitv in the brightness as recorded in the Strouuugren b and y bands. or about twice the ~0.1% variation in total irracliance that has been measured thus aa for the coutemporary Sun curing the solar cycle.," Inspection of the results of long-term, high precision photometry of solar-type stars by \citet[][their Fig. 17]{Lockwood97}237 as a function of the parameter $^{\prime}_{HK}$ suggests that this level of activity would correspond to an annual mean level of rms brightness variations of roughly 0.002 mag, i.e. variability in the brightness as recorded in the Strömmgren $b$ and $y$ bands, or about twice the $\sim 0.1\%$ variation in total irradiance that has been measured thus far for the contemporary Sun during the solar cycle."238 Therefore. we suggest that 0.2% represents au upper iuit to the likely excursion of the solar Iuuninous output as a result of cuhanced levels of magnetic activity.," Therefore, we suggest that $\sim 0.2\%$ represents an upper limit to the likely excursion of the solar luminous output as a result of enhanced levels of magnetic activity."239 We can refine this estimate of the upper lint further x nofiug that the brightucss changes given by Lockwoodetal.(1997) were for the mean variation of the smu of he Strónuueren hb aud y bands., We can refine this estimate of the upper limit further by noting that the brightness changes given by \citet{Lockwood97} were for the mean variation of the sum of the Strömmgren $b$ and $y$ bands.240 Hence. the variation of he total irraciance must be less than what is observed in these visible spectral bands.," Hence, the variation of the total irradiance must be less than what is observed in these visible spectral bands."241 Radicketal.(1998) estimated (for simall variations) a factor for converting )etwoeen ao fractional change in bolometric fux iuto he corresponding magnitude difference iu (5|y)/2.," \citet{Radick98} estimated (for small variations) a factor for converting between a fractional change in bolometric flux into the corresponding magnitude difference in $(b +242y)/2$."243 Adopting their conversion factor of 1.39 aud the estimate of brightuess variations of nunae eiven above viclds an estimate of for the uppor limit for variations in the bolometric fux., Adopting their conversion factor of 1.39 and the estimate of brightness variations of mag given above yields an estimate of for the upper limit for variations in the bolometric flux.244 This is only slightly larecr thu the mean variation of iu the total solar irradiance observed durius the course of the solar cvele., This is only slightly larger than the mean variation of in the total solar irradiance observed during the course of the solar cycle.245 The relatively more rapid rotation of 81152 invites further consideration iu the context of aneular momentum evolution aud the determination of stellar ages based on rotation. known as evrochrouologv (Barnes2007).," The relatively more rapid rotation of S1452 invites further consideration in the context of angular momentum evolution and the determination of stellar ages based on rotation, known as ""gyrochronology"" \citep{Barnes07}."246 At its measured (projected) rotation velocity and assumuime a stellar radius close to solar. the vevro-age’ of Sanders 1152 is 1.0+0.2 CC (Barues2007).," At its measured (projected) rotation velocity and assuming a stellar radius close to solar, the “gyro-age” of Sanders 1452 is $1.0 \pm 0.2$ Gyr \citep{Barnes07}."247. This value is an upper Πιτ because we mncasure only the projected velocity ¢sin/.," This value is an upper limit because we measure only the projected velocity $v\,\sin{i}$."248 A more direct comparison can be obtained with the esins age correlation given by Pace&Pasquini(2001.their 9)..," A more direct comparison can be obtained with the $v\,\sin{i}$ –age correlation given by \citet[][their Fig.2499]{Pace04}."250 The interred age (upper Iit) of 51152 based ou the three possible power law fits adopted by Pace Pasquini is iu the rauge of 1.21.5 Cox., The inferred age (upper limit) of S1452 based on the three possible power law fits adopted by Pace Pasquini is in the range of 1.2–1.5 Gyr.251 In either approach. the rotatiou-based age estiuiate for 81152 is in vivid coutrast to the age ranee for M67 of 3.5LS Cyr (Yadav or that of the Sun. namely. GGsyr (e...Do- and the simular solar age inuplied for the slow rotators of our sample.," In either approach, the rotation-based age estimate for S1452 is in vivid contrast to the age range for M67 of 3.5–4.8 Gyr \citep{Yadav08} or that of the Sun, namely, Gyr \citep[e.g.,][]{Bonanno02, Baker05} and the similar solar age implied for the slow rotators of our sample."252 Given its ligher rotational velocity. 81152 therefore represents an alternative path for angular momentum," Given its higher rotational velocity, S1452 therefore represents an alternative path for angular momentum"253"where z(8) is Earth's column density as a function of angle and c, is the total v-matter crossection.",where $z(\theta)$ is Earth's column density as a function of angle and $\sigma_T$ is the total $\nu$ -matter crossection.254 We have calculated the expected number of events for a km? detector., We have calculated the expected number of events for a $^3$ detector.255 We have used CTEQS5 for the neutrino-matter cross-section (Laietal2000).., We have used CTEQ5 for the neutrino-matter cross-section \citep{2000EPJC...12..375L}.256 We follow the calculation by Stanev1991) of average muon range., We follow the calculation by \citep{1991PhRvD..44..3543L} of average muon range.257 The Earth column density is taken from the Preliminary Earth Reference Model (Dziewonski&Anderson1981)., The Earth column density is taken from the Preliminary Earth Reference Model \citep{1981PEPI...25..297D}.258". Using approximations A and B and the neutrino spectrum derived from the Band function, we have studied a GRB with a photon break energy e"".=300 keV, located at a redshift z=1 and with Lorentz bulk boost I'=300."," Using approximations A and B and the neutrino spectrum derived from the Band function, we have studied a GRB with a photon break energy $\epsilon^b_\gamma259= 300$ keV, located at a redshift z=1 and with Lorentz bulk boost $\Gamma=300$."260 We also set the spectral indices to a.=—2 and 3.=—1., We also set the spectral indices to $\alpha_\gamma=-2$ and $\beta_\gamma=-1$.261 We have normalized the neutrino fluence of all three spectra to the same (arbitrary) value as described in section 3.., We have normalized the neutrino fluence of all three spectra to the same (arbitrary) value as described in section \ref{sec:photopion}.262" For this GRB the effective break energy is €,=110 keV. The neutrino break energy for approximation A is 5.98x10! eV and for approximation B it's 1.63x10!"" eV. For all cases we have fixed the synchrotron energy break at 10!"" eV. Figure 2. shows the three neutrino spectra."," For this GRB the effective break energy is $\bar{\epsilon}_\gamma=110$ keV. The neutrino break energy for approximation A is $\times 10^{14}$ eV and for approximation B it's $\times26310^{15}$ eV. For all cases we have fixed the synchrotron energy break at $^{16}$ eV. Figure \ref{fig:NuFlux} shows the three neutrino spectra."264 Again it is clear that approximation A is inadequate because it overestimates the contribution of low energy neutrinos., Again it is clear that approximation A is inadequate because it overestimates the contribution of low energy neutrinos.265 Figure 3 the ratio approximation A to B of expected number of events as a function cos(4) for this example GRB., Figure \ref{fig:ratio} the ratio approximation A to B of expected number of events as a function $\cos(\theta)$ for this example GRB.266 For all GRB locations in the sky we see that approximation A overestimates the expected number of events., For all GRB locations in the sky we see that approximation A overestimates the expected number of events.267 For steeper angles the ratio 1s larger. because for approximation B the characteristic neutrino energy is higher and therefore Earth's attenuation is higher.," For steeper angles the ratio is larger, because for approximation B the characteristic neutrino energy is higher and therefore Earth's attenuation is higher."268 We have shown that the usual choice to describe the photon spectra in the calculation of neutrino fluxes from GRBs overestimates the contribution of high energy photons (and therefore contribution of low energy neutrinos is overestimated)., We have shown that the usual choice to describe the photon spectra in the calculation of neutrino fluxes from GRBs overestimates the contribution of high energy photons (and therefore contribution of low energy neutrinos is overestimated).269 This results in a higher exepected event rate by a factor of z 2 for all models that make this assumption., This results in a higher exepected event rate by a factor of $\approx$ 2 for all models that make this assumption.270 The actual value of the overestimation of the expected number of events depends on the matter column depth that neutrinos must cross through Earth in the direction of the GRB., The actual value of the overestimation of the expected number of events depends on the matter column depth that neutrinos must cross through Earth in the direction of the GRB.271 Also we have shown that the typical neutrino energy is zz 1012 eV. The characteristic energy of neutrinos is a factor of ¢ larger than the values obtained by Guettaetal(2004) and a factor of 10 than the average value used by Waxman&Bahcall(1997).., Also we have shown that the typical neutrino energy is $\approx$ $^{15}$ eV. The characteristic energy of neutrinos is a factor of $e$ larger than the values obtained by \citet{2004APh....20..429G} and a factor of 10 than the average value used by \citet{1997PhRvL..78.2292W}.272 For back of the envelope calculations we provide a new approximation to the Band function that is adequate for the computation of neutrino spectra., For back of the envelope calculations we provide a new approximation to the Band function that is adequate for the computation of neutrino spectra.273 Kashti&Waxman(2005) have discussed the effects of muon and pion energy losses leading, \citet{2005PhRvL..95..181101K} have discussed the effects of muon and pion energy losses leading274The standard. theory of big bang nucleosvnthesis (SBBN: og. Boeseaard Steigman 1985: Steigman 1980: Walker et al.,"The standard theory of big bang nucleosynthesis (SBBN; e.g., Boesgaard Steigman 1985; Steigman 1989; Walker et al."275" 1991) accurately predicts the primordial abundances of the light elements D. ο, tHe and  Li. as a function of the cosmic barvon density. nxQiA7. or. equivalently. of the barvon-to-photon ratio. η."," 1991) accurately predicts the primordial abundances of the light elements D, $^3$ He, $^4$He and $^7$ Li, as a function of the cosmic baryon density, $\rho_{\mathrm{b}} \propto 276 \Omega_{\mathrm{b}} \, h^2$, or, equivalently, of the baryon-to-photon ratio, $\eta$."277 Since the barvon density is the sole. parameter. in he SBBN. observations of D. He. Πο and SLi in astrophysical environments. not. vet alecteck by subsequent. stellar evolution oller a direct way to infer the barvon density of the universe ancl to assess whether the SBBN theory correetly describes the first three minutes of the hot early universe.," Since the baryon density is the sole parameter in the SBBN, observations of D, $^3$ He, $^4$ He and $^7$ Li in astrophysical environments not yet affected by subsequent stellar evolution offer a direct way to infer the baryon density of the universe and to assess whether the SBBN theory correctly describes the first three minutes of the hot early universe."278 Nowadavs. there is still disagreement on which is the value of 5 as inferred from limits on theprimordial deuterium abundance from high-redshift absorption systems (mostly Danipec Lyman à or Lyman limit absorbers: e... Carswell et al.," Nowadays, there is still disagreement on which is the value of $\eta$ as inferred from limits on theprimordial deuterium abundance from high-redshift absorption systems (mostly Damped Lyman $\alpha$ or Lyman limit absorbers; e.g., Carswell et al."279 1994: Songaila et al., 1994; Songaila et al.280 1994: Tytler. Fan Burles 1996: Rugers Hogan 1996: Songaila. Wanipler Cowie 1997: Webb et al.," 1994; Tytler, Fan Burles 1996; Rugers Hogan 1996; Songaila, Wampler Cowie 1997; Webb et al."281" 1997: 3urles ""Tytler 1998: Levshakov. Ixegel Takahara 1998: O'Moeara et al."," 1997; Burles Tytler 1998; Levshakov, Kegel Takahara 1998; O'Meara et al."282 2001: Pettini Bowen 2001: Levshakoy et al., 2001; Pettini Bowen 2001; Levshakov et al.283 2002): estimates of the primordial ‘He abundance from oxtragalactic regions (e.g.. Olive. Skillman Steigman 1997: Izotov et al.," 2002); estimates of the primordial $^4$ He abundance from extragalactic regions (e.g., Olive, Skillman Steigman 1997; Izotov et al."284" 1999: Peimbert. Peimbert Luriciana 2002: see also Pagel 2000 for a review) and estimates of the. pristine ""Li content in old. metal-poor (FefL] 1.5 dex). warm (dc5700 Ix) clwarl stars in the solar neighbourhood (e.g.. Spite Spite 1982: Spite. Maillard Spite 1984: Spite Spite 1986: Itebolo. Molaro Beckman 1988: Thorburn 1992. 1994: Bonifacio Alolaro 1997: Vauclai Charbonnel 1995. LOOS: Vhéeaado Vauclair Wl: Pinsonneault ct al."," 1999; Peimbert, Peimbert Luridiana 2002; see also Pagel 2000 for a review) and estimates of the pristine $^7$ Li content in old, metal-poor ([Fe/H] $\le -1.5$ dex), warm $T_{\mathrm{eff}} \ge 5700$ K) dwarf stars in the solar neighbourhood (e.g., Spite Spite 1982; Spite, Maillard Spite 1984; Spite Spite 1986; Rebolo, Molaro Beckman 1988; Thorburn 1992, 1994; Bonifacio Molaro 1997; Vauclair Charbonnel 1995, 1998; Théaado Vauclair 2001; Pinsonneault et al."285 1999. 2002).," 1999, 2002)."286 Measurements of Ηο in regions in the outer parts of the Galactic disc should provide values very close to the primordial one. owing to the slow evolution of the disc in these regions. but they appear to poorly constrain the 7 range (e.g.. Rood et al.," Measurements of $^3$ He/H in regions in the outer parts of the Galactic disc should provide values very close to the primordial one, owing to the slow evolution of the disc in these regions, but they appear to poorly constrain the $\eta$ range (e.g., Rood et al."287 1998: Dania. Hood Balser 2002).," 1998; Bania, Rood Balser 2002)."288 Data on cosmic microwave background (CM) anisotroples oovide an alternative. independent method [or constraining η].," Data on cosmic microwave background (CMB) anisotropies provide an alternative, independent method for constraining $\eta$."289 The first. release of results from. the(WALAP: )onnett et al., The first release of results from the; Bennett et al.290 2003: Spergel et al., 2003; Spergel et al.291 2003) makes the CALB the prime cosmic barvometer. owing to the high. precision.," 2003) makes the CMB the prime cosmic baryometer, owing to the high precision."292 A combination of data with other finer scale CMD experiments. (ACBAR lxuo et. al., A combination of data with other finer scale CMB experiments (ACBAR – Kuo et al.293 2002 and CDI Pearson οἱ al., 2002 – and CBI – Pearson et al.294 2002) ancl with astronomical measurements of the power spectrum (2dbk Galaxy Reelshilt Survey measurements Pereival et al., 2002) and with astronomical measurements of the power spectrum (2dF Galaxy Redshift Survey measurements – Percival et al.295 2001 0 and Lyman a forest data Croft et al., 2001 – and Lyman $\alpha$ forest data – Croft et al.296 2002: Cinedin. Llamilton 2002) gives Onh?=0.0224 40.0009. or. equivalently. qocsi=6.1n (Spergel et. al. ," 2002; Gnedin Hamilton 2002) gives $\Omega_{\mathrm{b}} \, h^2 = 0.0224 297 \pm 0.0009$ , or, equivalently, $\eta_{\mathrm{10, \, CMB}} = 298 6.1^{+0.3}_{-0.2}$ (Spergel et al. ,"299where mo—n 10755.By adopting the SBBN predictions and using this y value. one obtains," where $\eta_{\mathrm{10}} \equiv 10^{10} \, \eta$ .By adopting the SBBN predictions and using this $\eta$ value, one obtains"300maenification of the density for. 1000.<Al«50.000.,"magnification of the density for $1000<M<50,000$."301 In this range. the density function is better represented by the function .025/4.," In this range, the density function is better represented by the function $.025/M$."302 The map approximaes the time between crossings of the SOS by considering 1e motion of my to be Ixeplerian., The map approximates the time between crossings of the SOS by considering the motion of $m_3$ to be Keplerian.303 losteack of considering he distribution. of the. shadow duration in terms of the number of iterations of 4 we can instead consider the distrixition of shadow times. fay where A ds the number of iterations of the orbits for which it was shadow-able.," Instead of considering the distribution of the shadow duration in terms of the number of iterations of $\varphi$ we can instead consider the distribution of shadow times, $t_M$ where $M$ is the number of iterations of the orbits for which it was shadow-able."304 “Phe soid line in Figure 13. represents the probability density ο ‘shadow time for the numerical experiments., The solid line in Figure \ref{fig:probden2} represents the probability density of shadow time for the numerical experiments.305 Again. the data can best. be approximated bv the exponential densiv [function lor £=.0005.," Again, the data can best be approximated by the exponential density function for $\xi = .0005$."306 The results. found. here are in agreement. for small shadow durations. with previous results by Hayes.(2003) which showed that shadow: clurations for larger N-body πω... have an exponential distribution and can be thought of as a Poisson process.," The results found here are in agreement, for small shadow durations, with previous results by \cite{Hayes} which showed that shadow durations for larger $N$ -body systems have an exponential distribution and can be thought of as a Poisson process."307 The above results confirm. for short. lived orbits. previous investigations (Llaves2003) that showed numerical shadow durations. AM. for gravitating svstems follow a Poisson process with a exponential density. function.," The above results confirm, for short lived orbits, previous investigations \citep{Hayes} that showed numerical shadow durations, $M$, for gravitating systems follow a Poisson process with a exponential density function."308 “Phe result found in this study suggests for longer lived. orbits. the density function is better. approximated bv a function proportional to 1/A7.," The result found in this study suggests for longer lived orbits, the density function is better approximated by a function proportional to $1/M$."309 This may be because the population of longer lived orbits tends to be dominated by stable orbits. however this has not been investigated.," This may be because the population of longer lived orbits tends to be dominated by stable orbits, however this has not been investigated."310 In section. 77... areas of phase-space where the refinement procedure is more likely to fail are characterized.," In section \ref{sec:5}, areas of phase-space where the refinement procedure is more likely to fail are characterized."311 These areas are near escape boundaries where there is sullicient stretching of phase-space to cause the refinement procedure to [fail to converge (o0 a less noisy orbit., These areas are near escape boundaries where there is sufficient stretching of phase-space to cause the refinement procedure to fail to converge to a less noisy orbit.312 Interestingly this seems to be due to the growth of the variational equations over one time step., Interestingly this seems to be due to the growth of the variational equations over one time step.313 This does not rule out the failure of the the refinement procedure by the accumulative effect of the growth of the variational equation associated with large Lyapunoy exponents as discussed. by Zhu, This does not rule out the failure of the the refinement procedure by the accumulative effect of the growth of the variational equation associated with large Lyapunov exponents as discussed by \cite{Zhu}.314andHayes(2009).. 1n Figure 11. dt is demonstrated that as the volume of phase-space representing capture orbits decreases. it comes increasingly cdillieult to shadow capture orbits.," In Figure \ref{fig:ECCvsPERCENT} it is demonstrated that as the volume of phase-space representing capture orbits decreases, it becomes increasingly difficult to shadow capture orbits."315 This is à result of the distribution o‘failures of the refinement xocedure seen in Figure 10..., This is a result of the distribution of failures of the refinement procedure seen in Figure \ref{fig:Edist}.316 As the volume of phase-space associated with capture decreases. capture orbits get pushed up against the boundary OD wrere the one-step growth of he variational equations causes he refinement procedure to al.," As the volume of phase-space associated with capture decreases, capture orbits get pushed up against the boundary $\partial \mathcal{D}_0$ where the one-step growth of the variational equations causes the refinement procedure to fail."317 Finally. it was found that he shadow cistance for an orbit is proportional to the numx of iterations of the map (Figures 7 and S)).," Finally, it was found that the shadow distance for an orbit is proportional to the number of iterations of the map (Figures \ref{fig:TvsEPSILON} and \ref{fig:eps09}) )."318 Ht was noticed that if in addition to ἐν anc fy. orbits were required. to be shadow-able at the wll steps fye5 ancl yes. then initially shadow-able orbits continued to be shadow-able.," It was noticed that if in addition to $t_1$ and $E_1$, orbits were required to be shadow-able at the half steps $t_{1/2}$ and $E_{1/2}$, then initially shadow-able orbits continued to be shadow-able."319 When shadowing at the half step was required. the shadow distance typically increased »v about a [actor of two.," When shadowing at the half step was required, the shadow distance typically increased by about a factor of two."320 The Sitnikoy problem cliscussed in this study provides a straight forward wav of characterizing a domain of initial conditions as well as regions of stable and unstable motion., The Sitnikov problem discussed in this study provides a straight forward way of characterizing a domain of initial conditions as well as regions of stable and unstable motion.321 Work in progress considers slight changes to the Sitnikov problem in order to study shaclowing of unstable orbits., Work in progress considers slight changes to the Sitnikov problem in order to study shadowing of unstable orbits.322 For example. Soulisetal.(2007) consider slight perturbations to the mass ancl position (away from the z-axis) of m; and delineate regions of stable anc unstable motion.," For example, \cite{Soulis1} consider slight perturbations to the mass and position (away from the $z$ -axis) of $m_3$ and delineate regions of stable and unstable motion."323 Lt would be expected that. like the results found in this stucl. shadowing with the refinement procedure breaks down near boundaries of escape for unstable orbits.," It would be expected that, like the results found in this study, shadowing with the refinement procedure breaks down near boundaries of escape for unstable orbits."324 In fact. the break down of the refinement procedure near escape boundaries would. be expected. for general 8-bock configurations.," In fact, the break down of the refinement procedure near escape boundaries would be expected for general 3-body configurations."325 As solutions approach parabolic escape boundaries. an orbit can unclergo increasinglv long ejections from the left-over binary svstem.," As solutions approach parabolic escape boundaries, an orbit can undergo increasingly long ejections from the left-over binary system."326 Small changes in the energy of an orbit in this region can cause significant. changes in the time of return for the orbit., Small changes in the energy of an orbit in this region can cause significant changes in the time of return for the orbit.327 Lf the refinement procedure could make changes to the orbits so as to conserve the energy of the ejected body it mieht improve the success rate of the refinement. procedure., If the refinement procedure could make changes to the orbits so as to conserve the energy of the ejected body it might improve the success rate of the refinement procedure.328 Finally. the Sitnikov «ον problem (Souliset.al.2008) provides a starting point for cxamining the relationship between the shadowing distance and the number of bodies.," Finally, the Sitnikov 4-body problem \citep{Soulis2}329 provides a starting point for examining the relationship between the shadowing distance and the number of bodies."330(Swan bands with the strongest one having a bandhead at 5165 A)).,(Swan bands with the strongest one having a bandhead at 5165 ).331 Because of the large heliocentric distance CN is expected to be dominant. if some emission bands are apparent.," Because of the large heliocentric distance CN is expected to be dominant, if some emission bands are apparent."332 This is due to the fact that the ratio CN/C» increases with heliocentric distance (2) because CN Is less sensitive to the heliocentric distance., This is due to the fact that the ratio $_2$ increases with heliocentric distance \citep{ahearn:1981} because CN is less sensitive to the heliocentric distance.333 So far CN is also the only emission band detected in the visible range at large heltocentric distance (on Chiron. see ?)).," So far CN is also the only emission band detected in the visible range at large heliocentric distance (on Chiron, see \cite{bus:1991}) )."334contaminate it (more so at high energies. because of the energy dependence of the PSF). so that one cannot draw definite conclusions on the spectral energy distribution of this source.,"contaminate it (more so at high energies, because of the energy dependence of the PSF), so that one cannot draw definite conclusions on the spectral energy distribution of this source."335 ISO and ccolors for the unresolved system indicate a Flat Spectrum source., ISO and colors for the unresolved system indicate a Flat Spectrum source.336 In the X-ray the star shows significant variability. with variations of a factor of two on a 20 ks time scale.," In the X-ray the star shows significant variability, with variations of a factor of two on a 20 ks time scale."337 With KT=2.5 keV and N(H)=3.6x107 cm-. equivalent to Ay=19 mag. the source is relatively hot and deeply embedded in the cloud. [," With $kT =3382.5$ keV and $\nh = 3.6\times 10^{22}$ $^{-2}$ , equivalent to $A_V =33919$ mag, the source is relatively hot and deeply embedded in the cloud. ["340EC92] 95 is à heavily absorbed intermediate-mass (~ 4Ma) YSO. for which Preibisch(1999). derived. from IR spectroscopy. a spectral type K2. with Αν~36 mag and Ly~60L.,"EC92] 95 is a heavily absorbed intermediate-mass $\sim 4 M_{\sun}$ ) YSO, for which \cite{pre99} derived, from IR spectroscopy, a spectral type K2, with $A_V \sim 36$ mag and $L_{\rm bol} \sim34160\,L_\odot$."342 Assuming that the X-ray absorption was the same as the optical one. Preibisch(1998) attributed an extremely large X-ray luminosity to this star (Lx~ 109). much higher than any known stellar source.," Assuming that the X-ray absorption was the same as the optical one, \cite{pre98} attributed an extremely large X-ray luminosity to this star $L_{\rm X} \sim 10^{33}$ ), much higher than any known stellar source."343 oobservations (Preibisch.2003). allowed the X-ray absorption to be directly determined. showing it to be over a factor of two lower than the optical estimate. thus bringing the. X-ray luminosity to a value within the range observed in YSOs.," observations \citep{pre2003} allowed the X-ray absorption to be directly determined, showing it to be over a factor of two lower than the optical estimate, thus bringing the X-ray luminosity to a value within the range observed in YSOs."344 The reason why the two absorption determinations (optical and ray) differ is. however. still unknown.," The reason why the two absorption determinations (optical and X-ray) differ is, however, still unknown."345" The X-ray spectral parameters determined from this oobservation are KT=2.740.2 keV and Ly=1.9xI0?! ere s7!, with Ny=(3.2£0.1)«107 em and Z=0.2Zs. remarkably consistent with the parameters derived by Preibisch(2003) for the oobservation. showing that the source has little long-term variability in its X-ray emission,"," The X-ray spectral parameters determined from this observation are $kT=2.7\pm0.2$ keV and $L_{\rm X} = 1.9\times 10^{31}$ erg $^{-1}$, with $N_{\rm H} = (3.2\pm 0.1) \times 10^{22}$ $^{-2}$ and $Z = 0.2\,Z_\odot$, remarkably consistent with the parameters derived by \cite{pre2003} for the observation, showing that the source has little long-term variability in its X-ray emission."346 Preibisch(2003) speculated about à number of possible scenarios for the discrepancy. including some unidentified peculiarity of flat-spectrum YSOs. by analogy with IRS 5 in L1551.," \cite{pre2003} speculated about a number of possible scenarios for the discrepancy, including some unidentified peculiarity of flat-spectrum YSOs, by analogy with IRS 5 in L1551."347 However. the line of reasoningσ of Preibisch(2003) is based on the assumption that X-rayy emission is observed from the IRS 5 protostar (which has Ay-150 mag).," However, the line of reasoning of \cite{pre2003} is based on the assumption that X-ray emission is observed from the IRS 5 protostar (which has $A_V \sim 150$ mag)."348 As demonstrated by Favataetal.(2002) and Ballyetal.(2003) the X-ray emission in this case comes from the associated Herbig-Haro object HH 154. which has. given that the jet penetrates through the absorbing cloud. a much lower absorption. with Ay~7 mag. compatible with the absorption measured for the X-ray source.," As demonstrated by \cite{ffm+02} and \cite{bfr2003} the X-ray emission in this case comes from the associated Herbig-Haro object HH 154, which has, given that the jet penetrates through the absorbing cloud, a much lower absorption, with $A_V \sim 7$ mag, compatible with the absorption measured for the X-ray source."349 One possible explanation for the peculiar. discrepancy between the X-ray and optically determined absorption columns is that the source is surrounded by a thick acereting disk. and is seen nearly edge-on (explaining the large optically determined A).," One possible explanation for the peculiar discrepancy between the X-ray and optically determined absorption columns is that the source is surrounded by a thick accreting disk, and is seen nearly edge-on (explaining the large optically determined $A_V$ )."350 If at the same time the X-ray emission comes from a region somewhat displaced above or below the disk (as it could be the case for a large polar corona). this could explain the lower absorbing column density observed in X-rays.," If at the same time the X-ray emission comes from a region somewhat displaced above or below the disk (as it could be the case for a large polar corona), this could explain the lower absorbing column density observed in X-rays."351 However. this implies a corona with a significant scale height above the photosphere. and this ts in conflict with recent evidence. for example. with the observed occurrence of rotational modulation in YSOs in Orion (Flaccomioetal.2005)). which implys that most coronal material is located in compact structures.," However, this implies a corona with a significant scale height above the photosphere, and this is in conflict with recent evidence, for example, with the observed occurrence of rotational modulation in YSOs in Orion \citealp{fms+05A}) ), which implys that most coronal material is located in compact structures."352 While large magnetic loops have been observed. also in Orion YSOs (Favataetal..2005a)). the consistency of the X-ray emission from [EC92] 95 over a number of years appears to rule out the hypothesis that one may be seeing a large flaring loop extending at a large distance from the star's photosphere.," While large magnetic loops have been observed, also in Orion YSOs \citealp{ffr+2005}) ), the consistency of the X-ray emission from [EC92] 95 over a number of years appears to rule out the hypothesis that one may be seeing a large flaring loop extending at a large distance from the star's photosphere."353 Indeed. if the observed source variability reffig:lc). is due to a flare. then the timescales of such variability (<20 ks) would imply a relatively compact structure (Lioop< δω).," Indeed, if the observed source variability \\ref{fig:lc}) ), is due to a flare, then the timescales of such variability $\la 20$ ks) would imply a relatively compact structure $L_{\rm loop} \la 5 R_{\sun}$ )."354 Another possible explanation for the discrepancy between the X-ray and optically determined absorption columns would be the presence of a lower nass companion., Another possible explanation for the discrepancy between the X-ray and optically determined absorption columns would be the presence of a lower mass companion.355 This companion star would have to be displaced from the disk of the main star. and be the main source of X-ray luminosity while contributing only a small fraction of the optical flux.," This companion star would have to be displaced from the disk of the main star, and be the main source of X-ray luminosity while contributing only a small fraction of the optical flux."356 evertheless. for such a companion star to have an X-ray luminosity of 1.9x10?!s7!.. its mass should be around 2—3M. as can be inferred by looking at the seatter plot of X-ray luminosity versus star mass in e.g. Flaccomioetal.(2003a) and therefore would provide significant contribution to the luminosity of the system.," Nevertheless, for such a companion star to have an X-ray luminosity of $1.9 \times 10^{31}$, its mass should be around $2-3357M_{\sun}$, as can be inferred by looking at the scatter plot of X-ray luminosity versus star mass in e.g. \citet{fdm03} and therefore would provide significant contribution to the luminosity of the system."358 This source lies at 3.7 arcesec from the position of the Herbig- flow 4456 identified by Davisetal.(1999)., This source lies at 3.7 arcsec from the position of the Herbig-Haro flow 456 identified by \citet{dmr+99}.359. We believe it unlikely that the X-ray emission originates. from shocks generated by the Herbig-Haro flow (as reported by citetpfe+2001 for 22 or Favataetal.(2002) for 1154)., We believe it unlikely that the X-ray emission originates from shocks generated by the Herbig-Haro flow (as reported by \\citet{pfg+2001} for 2 or \citet{ffm+02} for 154).360 The source is weak so its spectral characteristics cannot be well constrained. however it appears to be harder than generally observed in Herbig-Haro flows.," The source is weak so its spectral characteristics cannot be well constrained, however it appears to be harder than generally observed in Herbig-Haro flows."361 The source has à 2MASS counterpart at only 0.2 aresec distance and a corresponding. source classified às a Class HI star., The source has a 2MASS counterpart at only 0.2 arcsec distance and a corresponding source classified as a Class III star.362 A spectral fit withan absorbed IT plasma gives N(H)=(3.5€1.3)x107 em? and kT=4.9€4.5 (P= 0.85). consistent with the X-ray emission from the stellar corona of a CTTS or WTTS star.," A spectral fit withan absorbed 1T plasma gives $N({\rm H}) = (3.5 \pm 1.3) \times36310^{22}$ $^{-2}$ and $kT = 4.9 \pm 4.5$ $P = 0.85$ ), consistent with the X-ray emission from the stellar corona of a CTTS or WTTS star."364 This source has no 2MASS nor optical counterpart., This source has no 2MASS nor optical counterpart.365 It has no counterpart within the [SO catalogue by Kaasetal.(2004) and it was detected by, It has no counterpart within the ISO catalogue by \citet{kob+04} and it was detected by366the ring until the leading edge is reached again.,the ring until the leading edge is reached again.367 The whole pattern rotates around the ring with the period of the binary., The whole pattern rotates around the ring with the period of the binary.368 The lower panel of Fig.3 shows a particular case where the leading edge is located at 65220 and the mean decay angle governing the exponential 15 0.32 of 27 radians., The lower panel of Fig.3 shows a particular case where the leading edge is located at $\phi_0$ $^{\circ}$ and the mean decay angle governing the exponential is 0.32 of $\pi$ radians.369 The underlying physical picture is of some hotspot in the cireumbinary disk material. rotating with the binary and continually refreshing the glowing material in the ring.," The underlying physical picture is of some hotspot in the circumbinary disk material, rotating with the binary and continually refreshing the glowing material in the ring."370 For the example shown in the lower panel of Fig.3 the decay time for emission is thus about 4 days. (, For the example shown in the lower panel of Fig.3 the decay time for emission is thus about 4 days. (371This is appropriate for He I: for Πα the decay time is about 14 days.),This is appropriate for He I; for $\alpha$ the decay time is about 14 days.)372 This model. having as it does a sharp leading edge. is too simple to match reality. but the final step is to convolve the distribution given by (2) with a Gaussian function.," This model, having as it does a sharp leading edge, is too simple to match reality, but the final step is to convolve the distribution given by (2) with a Gaussian function."373 The resolution for these observations is about 0.2 in terms of the relative Doppler shift and the sharp edges in Fig.3 are removed., The resolution for these observations is about 0.2 in terms of the relative Doppler shift and the sharp edges in Fig.3 are removed.374 The effect of a Gaussian convolution ts shown in Fig.4. which shows the convolution of the lower panel of Fig.3 with a Gaussian of standard deviation 0.4.," The effect of a Gaussian convolution is shown in Fig.4, which shows the convolution of the lower panel of Fig.3 with a Gaussian of standard deviation 0.4."375 Note that the relative Doppler shift is shown between -2 and +2., Note that the relative Doppler shift is shown between -2 and +2.376 The standard deviation of 0.4 ts rather larger than resolution. but is supposed to contain some spread in the orbital speeds of the inner rim of the circumbinary disk.," The standard deviation of 0.4 is rather larger than resolution, but is supposed to contain some spread in the orbital speeds of the inner rim of the circumbinary disk."377 It has not been finely tuned but was chosen so that the ratio of separation of the two peaks in He to the FWHM be about right., It has not been finely tuned but was chosen so that the ratio of separation of the two peaks in $\alpha$ to the FWHM be about right.378 The model was used to calculate the form of the observations shown in Fig.1 (and later figures) in the following way., The model was used to calculate the form of the observations shown in Fig.1 (and later figures) in the following way.379 The predicted convolved spectrum was calculated in terms of the relative Doppler shift for a number of different values of the decay constant and for every 20° in do., The predicted convolved spectrum was calculated in terms of the relative Doppler shift for a number of different values of the decay constant and for every $^{\circ}$ in $\phi_0$.380 The convolved spectra were ther fitted with two Gaussians (the amounts. positions and standard deviations all being free parameters): very good representations were obtained in all cases.," The convolved spectra were then fitted with two Gaussians (the amounts, positions and standard deviations all being free parameters); very good representations were obtained in all cases."381 In Fig.4 is shown not only the convolved spectrum but also the two fitted Gaussians and their sum., In Fig.4 is shown not only the convolved spectrum but also the two fitted Gaussians and their sum.382 It is notable that in this case the centres of both fitted Gaussians are on the blue shifted side: this example corresponds to JD +246.5 and 4259.5 in Figs.l and 2., It is notable that in this case the centres of both fitted Gaussians are on the blue shifted side; this example corresponds to JD +246.5 and +259.5 in Figs.1 and 2.383 The results of these calculations are shown in Fig.2 and may be directly compared with the data in Fig.J., The results of these calculations are shown in Fig.2 and may be directly compared with the data in Fig.1.384 Figure 2 shows the calculated speeds for the red and blue components. after scaling to an assumed orbital speed for the ring of 250 km sl. with an assumed systemic velocity of 70 km s! and phase oo equal to zero on JD —-245.5.," Figure 2 shows the calculated speeds for the red and blue components, after scaling to an assumed orbital speed for the ring of 250 km $^{-1}$, with an assumed systemic velocity of 70 km $^{-1}$ and phase $\phi_0$ equal to zero on JD +245.5."385 The same parameters are used in Figs., The same parameters are used in Figs.386 5 and 6., 5 and 6.387 For the H« data the railroad. track. behavior of the two narrow components invited calculation of the rotational speed in terms of half the difference of the red and blue velocities and the apparent systemic recessional speed from the mean., For the $\alpha$ data the railroad track behavior of the two narrow components invited calculation of the rotational speed in terms of half the difference of the red and blue velocities and the apparent systemic recessional speed from the mean.388 Both quantities were rather constant with time: Fig., Both quantities were rather constant with time; Fig.389 3 of Blundell. Bowler & Schmidtobreick (2008). reproduced in Fig.7 of this paper.," 3 of Blundell, Bowler $\&$ Schmidtobreick (2008), reproduced in Fig.7 of this paper."390 I have displayed the He I data in the same way. Figs.," I have displayed the He I data in the same way, Figs."391 5 and 6., 5 and 6.392 Fig.5 shows the nominal rotational speed (half the difference) of the cireumbinary disk obtained from He I 6678 aand from the model as explained above., Fig.5 shows the nominal rotational speed (half the difference) of the circumbinary disk obtained from He I 6678 and from the model as explained above.393 The agreement between the data and the model is very good., The agreement between the data and the model is very good.394 In the model, In the model395Tot Jupiters are Jupiter-mass planets located within  0.1 AU or less from their parent stars.,Hot Jupiters are Jupiter-mass planets located within $\sim$ 0.1 AU or less from their parent stars.396 Because of (he close proximity to (he parent stars. hot Jupiters have been expected to be able to influence their stellar companions via magnetic Schaeler2000) and/or tidal interactions (Linet 2008).," Because of the close proximity to the parent stars, hot Jupiters have been expected to be able to influence their stellar companions via magnetic \citep{Cuntz,RS00} and/or tidal interactions \citep{Lin96,Jackson09,Pfahl08}."397. The observations of Ca II 11 Ix lines from a number of stus harboring a hot Jupiter have suggested that the chromospheric activities. characterized either by line intensity. or short-time variability. sometimes correlate with the orbits of their planets 2003.2005. 2008).," The observations of Ca II H K lines from a number of stars harboring a hot Jupiter have suggested that the chromospheric activities, characterized either by line intensity or short-time variability, sometimes correlate with the orbits of their planets \citep{Shk03,Shk05,Shk08}."398. These phenomena can be modelled as a hot spot or a more “variable” region. despite residing in the chromosphere. following the planets orbital motion with a phase dillerence.," These phenomena can be modelled as a hot spot or a more “variable"" region, despite residing in the chromosphere, following the planet's orbital motion with a phase difference."399" In. particular. the observations carried out in 2001. 2002. and 2005 imply that a hot spot on HD 179949 persistently leads the planet by 70 with the intensity of ~10%"" erg/s in Ca LL emissions."," In particular, the observations carried out in 2001, 2002, and 2005 imply that a hot spot on HD 179949 persistently leads the planet by $\sim 70^\circ$ with the intensity of $\sim 10^{27}$ erg/s in Ca II emissions."400 The similar phase lead of a variable region in optical has been suggested by the MOST satellite photometry for the hot-Jupiter host star 7. Dootis (Walkerοἱal.2003)., The similar phase lead of a variable region in optical has been suggested by the MOST satellite photometry for the hot-Jupiter host star $\tau$ Bootis \citep{Walker08}.401". Since the planet-induced stellar activities occur only. once during one orbital period. the origins of this ""spot have been attributed to magnetic rather than tidal inleractions."," Since the planet-induced stellar activities occur only once during one orbital period, the origins of this “spot"" have been attributed to magnetic rather than tidal interactions."402 One of the commonly adopted scenarios to describe (he star-planet magnetic interactions is (he magnetic interactions between Jupiter and its Galilean satellites (see Zarka 2007 [or a review)., One of the commonly adopted scenarios to describe the star-planet magnetic interactions is the magnetic interactions between Jupiter and its Galilean satellites (see Zarka 2007 for a review).403 In (his scenario. the orbital motion of the Galilean satellites relative to Jupiter's nagnetosphere taps the orbital energv of the satellites αἱ a rate (hat depends on detailed nodelling on the magnetic interactions.," In this scenario, the orbital motion of the Galilean satellites relative to Jupiter's magnetosphere taps the orbital energy of the satellites at a rate that depends on detailed modelling on the magnetic interactions."404 The interactions can be classified into (wo types: the unipolar interaction wilh an unmaegnetized satellite such as Io and the magnetic reconnection with a magnetized satellite such as Ganvmede., The interactions can be classified into two types: the unipolar interaction with an unmagnetized satellite such as Io and the magnetic reconnection with a magnetized satellite such as Ganymede.405 The energy is then transported by the Alfvénn waves and/or bv a fast electron beam along the field lines from the satellites to Jupiter's surface where (he energv is dissipated. thereby. explaining the satellite-inducecl emissions from Jupiter.," The energy is then transported by the Alfvénn waves and/or by a fast electron beam along the field lines from the satellites to Jupiter's surface where the energy is dissipated, thereby explaining the satellite-induced emissions from Jupiter."406 In the case of star-planet interactions. this picture has been modified to take into consideration stellar winds along; open field lines or to allow lor a large static magnetic loop connecting the star and the planet.," In the case of star-planet interactions, this picture has been modified to take into consideration stellar winds along open field lines or to allow for a large static magnetic loop connecting the star and the planet."407 As a result. the phase clifferences between the stellar “spot” and (he planet are explained by the time lag due to the Allvénn travel time io the star in the Allveen-wave model (Preusseetal.2006).. or bv the large magnetic loop having a geometry across longitudes of the star in the electron-beam case (Lanza2008).," As a result, the phase differences between the stellar “spot"" and the planet are explained by the time lag due to the Alfvénn travel time to the star in the Alfv́een-wave model \citep{Pre06}, or by the large magnetic loop having a geometry across longitudes of the star in the electron-beam case \citep{Lanza08}."408.. A (hree-climensional resistive magnetohyvdrodynamic simulation was performed to study how the magnetic field-aligned current can develop from a hot Jupiter (Preusseetal.2007)., A three-dimensional resistive magnetohydrodynamic simulation was performed to study how the magnetic field-aligned current can develop from a hot Jupiter \citep{Pre07}.409. llowever. how the stellar atinosphere thermally responds (o any energy injection from the planet so as to generate the chromospheric. emissions. of. 2107*OF erg/s remains− elusive.," However, how the stellar atmosphere thermally responds to any energy injection from the planet so as to generate the chromospheric emissions of $\gtrsim 10^{27}$ erg/s remains elusive."410− :The, The411star counts in à ὁ respectively as shown in Fig.,star counts in $\alpha$ $\delta$ respectively as shown in Fig.412 3., 3.413 The differences between our estimated centre and the obtained one of Webda are shown in the figure., The differences between our estimated centre and the obtained one of Webda are shown in the figure.414 To establish the radial densitv profile (RDP) of Ru 15. we counted the stars within concentric shells in equal incremental steps (r<1) aremin [from (he cluster centre.," To establish the radial density profile (RDP) of Ru 15, we counted the stars within concentric shells in equal incremental steps $r\leq1$ ) arcmin from the cluster centre."415 We repeated (his process for 1<rx2 up tor <I0 avemin. ie. the stellar density is derived out to the preliminary radius of the cluster.," We repeated this process for $1<r\leq2$ up to $r\leq10$ arcmin, i.e. the stellar density is derived out to the preliminary radius of the cluster."416 The stars of the next steps should be subtracted from the later ones. so that we obtained only the amount of the stars within the relevant shell's area. not a cumulative count.," The stars of the next steps should be subtracted from the later ones, so that we obtained only the amount of the stars within the relevant shell's area, not a cumulative count."417 Finally. we divided the star counts in each shell to the area of that shell those stars belong to.," Finally, we divided the star counts in each shell to the area of that shell those stars belong to."418 The density uncertainties in each shell was calculated using Poisson noise statistics., The density uncertainties in each shell was calculated using Poisson noise statistics.419 Fig., Fig.420 4 shows the RDP from the new centre of Ru 15 to the maximum angular separation of 5 arcmin where the stability of density has been reached., 4 shows the RDP from the new centre of Ru 15 to the maximum angular separation of 5 arcmin where the stability of density has been reached.421 To determine the structural parameters of (he cluster more precisely. we applied the empirical King model (1966).," To determine the structural parameters of the cluster more precisely, we applied the empirical King model (1966)."422" The Ning model parameterizes the density function p(r) as: where /j,. fy aud r. ave background. central star density and the core radius of the cluster respectively."," The King model parameterizes the density function $\rho(r)$ as: where $f_{bg}$, $f_{0}$ and $r_{c}$ are background, central star density and the core radius of the cluster respectively."423 From (he concentration parameter c. defined as ο=ηD). Nilakshi et al. (," From the concentration parameter $c$, defined as $c= (R_{lim}/R_{core})$, Nilakshi et al. ("4242002) concluded that the angular size of the coronal region is about 6 times the core radius.,2002) concluded that the angular size of the coronal region is about 6 times the core radius.425" Maciejewski Niedzielski (2007) reported Chat πρ may vary for individual clusters between about 222... and 74,44.", Maciejewski Niedzielski (2007) reported that $R_{lim}$ may vary for individual clusters between about $2 R_{core}$ and $7 R_{core}$.426 In our case. we can see that ΠΟ οτι. ie. it lies within the previous vales.," In our case, we can see that $R_{lim}$ $6.9 R_{core}$, i.e. it lies within the previous vales."427 The cluster limited radius can be defined at that radius which covers (he entire cluster area and reaches enough stability with the background. density. ie. the difference between the observed density profile and the background one is almost equal zero.," The cluster limited radius can be defined at that radius which covers the entire cluster area and reaches enough stability with the background density, i.e. the difference between the observed density profile and the background one is almost equal zero."428 It is noted that the determination of a cluster radius is made by. (he spatial coverage and unilormity of PPMXL photometry which allows one to obtain reliable data on the projected distribution of stars for large extensions to the clusters’ halos., It is noted that the determination of a cluster radius is made by the spatial coverage and uniformity of PPMXL photometry which allows one to obtain reliable data on the projected distribution of stars for large extensions to the clusters' halos.429 On the other hand. the concentration parameter seems to be related to cluster age. i.e. lor clusters vounger than about 1 Gyr. it tends (ο increase with eluster age.," On the other hand, the concentration parameter seems to be related to cluster age, i.e. for clusters younger than about 1 Gyr, it tends to increase with cluster age."430 Nilakshi et al. (, Nilakshi et al. (4312002) notes that the halos’ sizes are smaller for olcler svstems.,2002) notes that the halos' sizes are smaller for older systems.432 Finally. we can infer that open clusters appear to be somewhat larger in (he near-infrared than in the optical data. Sharma et al. (," Finally, we can infer that open clusters appear to be somewhat larger in the near-infrared than in the optical data, Sharma et al. ("4332006).,2006).434 Knowing the clusters total mass (Sec., Knowing the cluster's total mass (Sec.435 3.3). the tidal radius can be given by applying the equation of Jeffries οἱ al. (," 3.3), the tidal radius can be given by applying the equation of Jeffries et al. ("436"2001): where A) and Af, ave (he tidal radius and total mass of the cluster respectively.",2001): where $R_{t}$ and $M_{c}$ are the tidal radius and total mass of the cluster respectively.437"such as 41335 and 41240 (--600+50kms!,Hillieretal.2001).","such as $\lambda$ 1335 and $\lambda$ 1240 \citep[$-600 \pm 50 ~\kms$,][]{hillier01}."438". This range of velocities is consistent withHe1 A10833 being formed in the wind of Eta Car A. Eta Car B might significantly influence the 410833 line profile through photoionization of the outer parts of the wind of Eta Car A and.> in. this. case, the amount of emission∙∙ and absorption∙ seen in∙Het A10833 will strongly depend on the orbital of the system."," This range of velocities is consistent with $\lambda$ 10833 being formed in the wind of Eta Car A. Eta Car B might significantly influence the $\lambda$ 10833 line profile through photoionization of the outer parts of the wind of Eta Car A and, in this case, the amount of emission and absorption seen in $\lambda$ 10833 will strongly depend on the orbital parameters of the system."439" Nevertheless, at phases sufficiently farparameters from periastron such as at @=11.875, there is no evidence for additional velocity fields in our line-of-sight, such as one would expect in the case that the absorption was formed in the wind of Eta Car B or in high-velocity material from the wind-wind collision zone."," Nevertheless, at phases sufficiently far from periastron such as at $\phi=11.875$, there is no evidence for additional velocity fields in our line-of-sight, such as one would expect in the case that the absorption was formed in the wind of Eta Car B or in high-velocity material from the wind-wind collision zone."440The profiles from ¢=11.991 and ó=11.998 show very different 410833 absorptions compared to theone from à=11.875.,The profiles from $\phi=11.991$ and $\phi=11.998$ show very different $\lambda$ 10833 absorptions compared to theone from $\phi=11.875$.441" The low-velocity: absorption: strengthened, becoming. nearly saturated from -40 to kms™,−↥∙ with the exception of the equatorial ejecta emission-580 at —250kms""."," The low-velocity absorption strengthened, becoming nearly saturated from $-40$ to $-580~\kms$, with the exception of the equatorial ejecta emission at $-250~\kms$."442" A broad, high-velocity absorption ranging from —580 to —-1900kms""! appeared by ϕ=11.991 and strengthened by $=11.998."," A broad, high-velocity absorption ranging from $-580$ to $-1900~\kms$ appeared by $\phi=11.991$ and strengthened by $\phi=11.998$."443" We would not expect this high-velocity absorption from the velocity field of the wind of Eta Car A, as seen in the spectrum from ¢=11.875."," We would not expect this high-velocity absorption from the velocity field of the wind of Eta Car A, as seen in the spectrum from $\phi=11.875$."444" Therefore, the 210833 absorption line profile strongly indicates that, in addition to the wind of Eta Car A, at least one more velocity structure is crossing our line-of-sight to Eta Car."," Therefore, the $\lambda$ 10833 absorption line profile strongly indicates that, in addition to the wind of Eta Car A, at least one more velocity structure is crossing our line-of-sight to Eta Car."445" That high-velocity absorption is transient; as the considerablspectrum observed at ¢=12.014 hows, it kms!has faded y and isΡpresent only urup to li-900900kms."," That high-velocity absorption is transient; as the spectrum observed at $\phi=12.014$ shows, it has faded considerably and ispresent only up to $-900~\kms$."446" By ὁ°=12.041, theHer 410833 profile ist quite similar to that recorded at ὁ=11.875, but the —40 to —580kms! absorption is saturated, indicating a higher column density of Het."," By $\phi=12.041$, the $\lambda$ 10833 profile is quite similar to that recorded at $\phi=11.875$, but the $-40$ to $-580~\kms$ absorption is saturated, indicating a higher column density of ."447" The 410833 and 420587 line profiles, recorded at ¢=11.998, demonstrate that the high-velocity absorption component is much stronger in 410833 than in the 420587 line (Fig. 2))."," The $\lambda$ 10833 and $\lambda$ 20587 line profiles, recorded at $\phi=11.998$, demonstrate that the high-velocity absorption component is much stronger in $\lambda$ 10833 than in the $\lambda$ 20587 line (Fig. \ref{fig2}) )."448" While noticeable from —600 to —1000kms-!, the absorption is 1much weaker inthe range of —1100 to —1600kms-'. ! The 410833 line -2p element[][3]P)) absorption— originates⋅⋅ from the metastable triplet state, while the Her 420587 line — 1]P))originates from the metastable state."," While noticeable from $-600$ to $-1000~\kms$, the absorption is much weaker inthe range of $-1100$ to $-1600~\kms$ The $\lambda$ 10833 line – ) absorption originates from the metastable triplet state, while the $\lambda$ 20587 line – )originates from the metastable state."449 The population of the, The population of the450"smoother. xexp[7(v/1,)""C7] (Fritz1989:Zirakashvili&Aharonian 2007)).","smoother, $\propto \exp[-(\nu/\nu_{*})^{\beta/(\beta+2)}]$ \citealt{frit89,zirak07}) )."451" The position of the svuchrotvon peak flux. νο. is then also dependent on 2. and one can show that for 3=1 (or a,=0 in the previous notation) an important [actor ~10 arises. so that ij=9.51. whereas for 2=3 the svuchrotron peak corresponds approximately to the electron cut-off as py=1.2/5. (e.g.. Fig. 3))."," The position of the synchrotron peak flux, $\nu_p$, is then also dependent on $\beta$, and one can show that for $\beta=1$ (or $\alpha_p=0$ in the previous notation) an important factor $\sim 10$ arises, so that $\nu_{p}=9.5\nu_{c}$, whereas for $\beta=3$ the synchrotron peak corresponds approximately to the electron cut-off as $\nu_{p}=1.2\nu_{c}$ (e.g., Fig. \ref{SSCmax}) )."452 Expansion of the source could change the conclusions drawn above., Expansion of the source could change the conclusions drawn above.453 In particular. if one assiiies a verv low magnetic field such that svnchrotron losses are negligible. then adiabatie losses may become important and alter the electron distribution.," In particular, if one assumes a very low magnetic field such that synchrotron losses are negligible, then adiabatic losses may become important and alter the electron distribution."454 In this section. we examine the behavior of the svstem for a power-law electron distribution wilh a hieh value of the cut-off discussed above.," In this section, we examine the behavior of the system for a power-law electron distribution with a high value of the low-energy cut-off discussed above."455 For simplicitv. we consider a spherical source that expands with a constant. velocity à. The relativistic electron population will be affected by svinchirotron losses. and by aciabatic losses (e.g.. Longair 1932). As the emission region expands. the magnetic field decreases.," For simplicity, we consider a spherical source that expands with a constant velocity $u$, The relativistic electron population will be affected by synchrotron losses, and by adiabatic losses (e.g., Longair 1982), As the emission region expands, the magnetic field decreases."456" We consider a scaling (L/D0"" with 1<onx2 to study the evolution of the svstem."," We consider a scaling $B \propto (1/R)^m457\propto (1/t)^m$ with $1\leq m\leq 2$ to study the evolution of the system."458 The limiting value mm=2 corresponds (o conservation of magnetic flux for the longitudinal component. whereas m.—1 holds for the perpendicular component. (," The limiting value $m=2$ corresponds to conservation of magnetic flux for the longitudinal component, whereas $m=1$ holds for the perpendicular component. ("459Note that [ον m=1 the ratio of the electrons’ energy density to (he magnetic field energy density remains constant).,Note that for $m=1$ the ratio of the electrons' energy density to the magnetic field energy density remains constant).460 Which enerev loss process then determines (he electron behavior depends mainly on (he magnetic field strength ancl the size of the source., Which energy loss process then determines the electron behavior depends mainly on the magnetic field strength and the size of the source.461 A simple comparison of the above relations shows that when Z4> P. Le. aciabatic losses dominate over radiative losses.," A simple comparison of the above relations shows that when $P_{\rm ad}>P_{\rm syn}$ , i.e., adiabatic losses dominate over radiative losses."462 For example. if one considers expansion al speed 4~e and an initial source dimension2y~LO! em. then [or energies below," For example, if one considers expansion at speed $u \sim c$ and an initial source dimension$R_0\sim 10^{14}$ cm, then for energies below"463Paper I showed that the region of maximun polarization is offset by oA+35° along the ecliptic plane Irom the laree-erain inflow direction.,Paper I showed that the region of maximum polarization is offset by $ \delta \lambda \sim +35\deeg$ along the ecliptic plane from the large-grain inflow direction.464 The present note provides support lor this hypothesis (hat the polarization originates wilh interstellar dust at the heliosphere. using additional polarization data (§??)). evidence for olivine grains wilh stable alignment as the grains approach the heliosphere (§8??.. 77)). and a diseussion of the still-uncertain grain alignment mechanisms in the context of the heliosphere interaction with the ISM 7?7)).," The present note provides support for this hypothesis that the polarization originates with interstellar dust at the heliosphere, using additional polarization data \ref{sec:data}) ), evidence for olivine grains with stable alignment as the grains approach the heliosphere \ref{sec:licdust}, \ref{sec:disrupt}) ), and a discussion of the still-uncertain grain alignment mechanisms in the context of the heliosphere interaction with the ISM \ref{sec:alignment}) )."465 Charged ISDGs spin rapidly and will always be aligned so that (he observed polarization.P.. of starlight is parallel to the magnetic field direction.hy. regardless of the alignment mechanism)mainBodyCitationEnd524]Lazarian:2003.," Charged ISDGs spin rapidly and will always be aligned so that the observed polarization, of starlight is parallel to the magnetic field direction, regardless of the alignment mechanism."466 Evidently aligned grains do not behave like dumb compass needles. but rather trace the coupling between (he erain angular momentum andByy.," Evidently aligned grains do not behave like dumb compass needles, but rather trace the coupling between the grain angular momentum and."467. Thus. the T32 data indicate that iis approximately parallel to the ealactic plane aud oriented towards (~90° ((see Fig.," Thus, the T82 data indicate that is approximately parallel to the galactic plane and oriented towards $\sim$ (see Fig."468 6 of T82). or as found here (~105°.," 6 of T82), or as found here $\sim$."469. A similar orientation for iis indicated by the 2.6 kIIz Langmuir emission events observed by VovagerIXGO3)., A similar orientation for is indicated by the 2.6 kHz Langmuir emission events observed by Voyager.470. Driangulation by Vovager 1 and 2 show that tlie dozen emission events detected in the 1990's arise in the outer heliosphere at ~100 AU from the Sun. and that these dozen enission events are approximately aligned with the galactic planemechanisms).," Triangulation by Voyager 1 and 2 show that the dozen emission events detected in the 1990's arise in the outer heliosphere at $\sim$ 100 AU from the Sun, and that these dozen emission events are approximately aligned with the galactic plane."471. The tenuous nature of the interstellar eloud surrounding the Sun reduces the collisional disruption of grain alignment, The tenuous nature of the interstellar cloud surrounding the Sun reduces the collisional disruption of grain alignment472The paramcters of the thick disk aud the power-law halo have wide 47 distributions.,The parameters of the thick disk and the power-law halo have wide $\chi^2$ distributions.473 The scale height of the thick disk (2.55 kpc) is about twice of that of the thin disk aud agrees with the 7;;=1 kpc (ih.=29/2) even by Burstein1979. for SO ealaxies., The scale height of the thick disk (2.55 kpc) is about twice of that of the thin disk and agrees with the $h_z=1$ kpc $h_z=z_0/2$ ) given by \cite{Burstein79} for S0 galaxies.474 The thick disk scale leneth is 11.03 kpc. larger than that of the thin disk.," The thick disk scale length is 11.03 kpc, larger than that of the thin disk."475 The Imminosity of the thick disk is of the thin disk., The luminosity of the thick disk is of the thin disk.476 The core radius of the power-law halo is 11.1 spe. iudicatiug a flat halo.," The core radius of the power-law halo is 14.4 kpc, indicating a flat halo."477 This agrees with the third component of three disk models of NJ and that of Shaw (μονο (1989). iu which the scale leneths of their third disks are about 13.9 kpe (at a distance of 11.5 Alpe).," This agrees with the third component of three disk models of NJ and that of Shaw Gilmore (1989), in which the scale lengths of their third disks are about 13.9 kpc (at a distance of 14.5 Mpc)."478 The values of + concentrate in the rauge of 3.2 to L0. which indicates that there does not exist ar7 halo in this galaxy.," The values of $\gamma$ concentrate in the range of 3.2 to 4.0, which indicates that there does not exist a $r^{-2}$ halo in this galaxy."479 The best power index is 3.88. between the 3.5 of Milkv Way (Zinn1985)) aud the of M31 (Pritchet&vandenBergh199 1)).," The best power index is 3.88, between the 3.5 of Milky Way \cite{Zinn85}) ) and the 4.0 of M31 \cite{PB94}) )."480" The (5,4; of the thin aud the thick disks slow several local minimus. owing to data sample."," The $r_{max}$ of the thin and the thick disks show several local minimums, owing to data sampling."481 The cutoffs of the thin aud thick disks are about 32 kpc aud 37 kpe. respectively.," The cutoffs of the thin and thick disks are about 32 kpc and 37 kpc, respectively."482 Analyses of correlation between anv two paraiucters show that some of the parameters are correlated. which means some of the parameters can not be determined indepeudenutlv.," Analyses of correlation between any two parameters show that some of the parameters are correlated, which means some of the parameters can not be determined independently."483 For example. we fluc correlation between tyy and zg» aud amoug pos. ry aud 5.," For example, we find correlation between $z_{01}$ and $z_{02}$ and among $\mu_{03}$, $r_0$ and $\gamma$."484 The similar «7 distiibutious of pgs. ry and 5 in Figure 8 also show such a correlation.," The similar $\chi^2$ distributions of $ \mu_{03}$, $ r_0$ and $ \gamma$ in Figure 8 also show such a correlation."485 The error for cach paramcter is obtained in the following way., The error for each parameter is obtained in the following way.486 Raucom values are selected for the observed data such that they obey a normal distribution. with sigias determined by the known errors iu each s;iuupled bin.," Random values are selected for the observed data such that they obey a normal distribution, with sigmas determined by the known errors in each sampled bin."487 We then obtain the bestfitted parameters for that set of data., We then obtain the best–fitted parameters for that set of data.488 This procedure is repeated three hundred times. giving us 300 separate determinations of the best-fitted value for cach parameter.," This procedure is repeated three hundred times, giving us 300 separate determinations of the best-fitted value for each parameter."489 The statistical standard deviation of each parameter from this procedure is adopted as the fal error for this paraiicter., The statistical standard deviation of each parameter from this procedure is adopted as the final error for this parameter.490 The best fit parameters and errors are listed in Table 6., The best fit parameters and errors are listed in Table 6.491 Table 6 also lists the total magnitude iu A-band to a surface brightuess of 28 mae 7. ug66n)=8.99 (or broad-band R = 9.10 frou the Zhou et al.," Table 6 also lists the total magnitude in $\rm\AA$ -band to a surface brightness of 28 mag $^{-2}$, $\rm m_{6660} = 8.99$ (or broad-band R = 9.10 from the Zhou et al."492 transformation). which is measured by replacing masked areas around the galaxy by the corrsponding parts of the muuasked galaxy.," transformation), which is measured by replacing masked areas around the galaxy by the corrsponding parts of the unmasked galaxy."493 Figure 5 presents the best fitting values for the three components compared to our data in both the « and directions., Figure 5 presents the best fitting values for the three components compared to our data in both the $z$ and $R$ directions.494 The fits are quite good in the range of R=10' to —δ for the z-profiles aud in the range of 2=30° to 1207 for R-profiles.," The fits are quite good in the range of $R = 1'$ to $R = 8'$ for the $z$ -profiles and in the range of $z= 30""$ to $120""$ for $R$ -profiles."495 As expected. the model deviates significantly from the data in the region of the bulge (not-Btted) aud where the warp of the disk becomes sizeable.," As expected, the model deviates significantly from the data in the region of the bulge (not-fitted) and where the warp of the disk becomes sizeable."496CAIDs of MA (NGC 6121) and. MIOT (NGC 6171) which have both BIID and RIID.,CMDs of M4 (NGC 6121) and M107 (NGC 6171) which have both BHB and RHB.497 Even ihe RIB of 47 Tue (NGC 104). whieh has only RIIDB. is clearly tilled.," Even the RHB of 47 Tuc (NGC 104), which has only RHB, is clearly tilted."498" So in the A, vs. (J—N,) CMDs it is difficult to determine the HB level of anv GGC and to correlate RGB Inuup positions and IIB levels in (he near-infrared CMDs. in contrast to the case for optical CMDSs."," So in the $K_{s}$ vs. $(J-K_{s})$ CMDs it is difficult to determine the HB level of any GGC and to correlate RGB bump positions and HB levels in the near-infrared CMDs, in contrast to the case for optical CMDs."499 second. the brightness interval between the MSTO (Alain Sequence Turnoff) and RGB lip is larger (9 mag) than in optical CMDs (76.5 mag). as can be easily seen in M4.," Second, the brightness interval between the MSTO (Main Sequence Turnoff) and RGB tip is larger $\sim$ 9 mag) than in optical CMDs $\sim$ 6.5 mag), as can be easily seen in M4."500" So the magnitude resolution in the A, vs. (/—νι) CAIDs is larger than in optical CMDs.", So the magnitude resolution in the $K_{s}$ vs. $(J-K_{s})$ CMDs is larger than in optical CMDs.501" However. color resolution in the A, vs. (J—Av.) CMDs is smaller (han in optical CALDs because in ihe latter the RGB and AGB (Asymptotic Giant Branch) are separated [rom one another at least in the lower part of the AGB."," However, color resolution in the $K_{s}$ vs. $(J-K_{s})$ CMDs is smaller than in optical CMDs because in the latter the RGB and AGB (Asymptotic Giant Branch) are separated from one another at least in the lower part of the AGB."502" In contrast. in the A, vs. (J—IN.) CMDs the RGB and AGB overlap. as clearly shown in the CMD of 47 Tue whose AGB is relatively rich."," In contrast, in the $K_{s}$ vs. $(J-K_{s})$ CMDs the RGB and AGB overlap, as clearly shown in the CMD of 47 Tuc whose AGB is relatively rich."503" In the optical CMDs of NGC 362. 47 Tuc. and M71 (NGC 6333). the separation of their RGBs and BIIDs is clearly seen but they partially overlap in the A, vs. (J—A.) CMDs."," In the optical CMDs of NGC 362, 47 Tuc, and M71 (NGC 6838), the separation of their RGBs and RHBs is clearly seen but they partially overlap in the $K_{s}$ vs. $(J-K_{s})$ CMDs."504" This also results from the fact Chat the color resolution in the A, vs. (J—dv.) CMDsS is lower than that in optical CAIDs.", This also results from the fact that the color resolution in the $K_{s}$ vs. $(J-K_{s})$ CMDs is lower than that in optical CMDs.505 Thircl. in the case of M22 (NGC 6656). its RGB is broader than those of other GGCs as found in the optical CMD (Peterson Cuclworth 1994).," Third, in the case of M22 (NGC 6656), its RGB is broader than those of other GGCs as found in the optical CMD (Peterson Cudworth 1994)."506 Moreover. in the lower part of the RGB. contamination from bulge component stars is severe.," Moreover, in the lower part of the RGB, contamination from bulge component stars is severe."507 To the right of the RGB of M22 there exists another RGB component., To the right of the RGB of M22 there exists another RGB component.508 This RGB component would be due to bulge RGB stars. which extend prominently reclwarcl in the optical CMD Gn our unpublished data) because of the strong blanketing effects of heavy metals as found in metal-rich. GGCs such as NGC 6553 (Ortolani. Darbuy. Dica 1990).," This RGB component would be due to bulge RGB stars, which extend prominently redward in the optical CMD (in our unpublished data) because of the strong blanketing effects of heavy metals as found in metal-rich GGCs such as NGC 6553 (Ortolani, Barbuy, Bica 1990)."509 Last. because NGC 362 and 47 Tuc CALDs are located in the direction of the SAIC. a contribution from the SAIC appears in (he lower right region of their CMDs.," Last, because NGC 362 and 47 Tuc CMDs are located in the direction of the SMC, a contribution from the SMC appears in the lower right region of their CMDs."510 In order to accurately measure (he luminosity of the RGB bump and construct luminosity funelions of the GGCs. we applied several standard procedures to delineate only the RGB sequences lor all GGCs except for M22 ancl MT1.," In order to accurately measure the luminosity of the RGB bump and construct luminosity functions of the GGCs, we applied several standard procedures to delineate only the RGB sequences for all GGCs except for M22 and M71."511 First we rejected. visually clear HB stars and AGB stars in the process of distinguishing outlving field stars [rom RGB stars., First we rejected visually clear HB stars and AGB stars in the process of distinguishing outlying field stars from RGB stars.512 Second. bv binning the RGB sequences in 0.5 mag intervals we measured (he average ancl sigma of each bin and rejected stars 2σ away from the mean value of each bin.," Second, by binning the RGB sequences in 0.5 mag intervals we measured the average and sigma of each bin and rejected stars $\sigma$ away from the mean value of each bin."513 We Chen remeasured, We then remeasured514IIul«t 1989: Dablem et al. 1995j).,Hulst \cite{HuHu}; Dahlem et al. \cite{DaLi}) ).515 However. the galaxy its been listed in the sample of TRAS bright ealaxies (Soifer ct al. 1987)).," However, the galaxy has been listed in the sample of IRAS bright galaxies (Soifer et al. \cite{So87}) ),"516 so there seenis clear evideuce for chhanced SE but it is presently not clear if 32011 josts a starburst nucleus., so there seems clear evidence for enhanced SF but it is presently not clear if 3044 hosts a starburst nucleus.517 In a spectroscopic study where he DIC. was investigated at two differcut slit positions oerpendicular to the galaxy disk. the positious of the detected DIC in the diagnostic diagrams fall in between he areas occupied by normal regions aud starburst (Tülliuaun Dettmar 2000)).," In a spectroscopic study where the DIG was investigated at two different slit positions perpendicular to the galaxy disk, the positions of the detected DIG in the diagnostic diagrams fall in between the areas occupied by normal regions and starburst (Tülllmann Dettmar \cite{TuDe}) )."518 Therefore no clear auswer of this debate can be eiven vet., Therefore no clear answer of this debate can be given yet.519 Even the term ‘starburst? is sometimes not clearly defined. as various researchers use different definitions.," Even the term `starburst' is sometimes not clearly defined, as various researchers use different definitions."520 We will come back to this point in 55, We will come back to this point in 5.521", The galaxy type is listed as SBc (Tully 1988)).", The galaxy type is listed as SBc (Tully \cite{Tu88}) ).522 Even classifications as a nonbarred galaxy (Sc) are listed (Nilson 1973))., Even classifications as a non–barred galaxy (Sc) are listed (Nilson \cite{Ni73}) ).523 There ave iudicatious that iun 33011a bar is present. aud kinematics by Lee biu (1997) has indeed discerued a bar.," There are indications that in 3044 a bar is present, and kinematics by Lee Irwin \cite{LeIr}) ) has indeed discerned a bar."524 It might be worth to notice that in the cighties a supernova (SNI1983E)) has been detected in this galaxy. (Darbon et al. 1989))., It might be worth to notice that in the eighties a supernova ) has been detected in this galaxy (Barbon et al. \cite{BaCa}) ).525 Iu Fig., In Fig.526 2 we present the Πα image., \ref{F2} we present the $\alpha$ image.527 The morphology of: the DIC shows various: ]features., The morphology of the DIG shows various features.528 An eDIG. laver can be detected at extraplanar distances up to z=0.8Ikpec.," An eDIG layer can be detected at extraplanar distances up to $z=0.8-1\,\rm{kpc}$."529 Several single plumes can also be diseerued., Several single plumes can also be discerned.530" South of the ealactic plane an extended structure is visible, which hasa looplike appearance."," South of the galactic plane an extended structure is visible, which has a loop–like appearance."531 This loop exteuds out to ~ L&kkpe with a radius of about Lkkpe. aud resembles the galactic πλ," This loop extends out to $\sim$ kpc with a radius of about kpc, and resembles the galactic supershells."532 The disk appears slightly warped. which is also apparent in the Roland image.," The disk appears slightly warped, which is also apparent in the R–band image."533 The Πα flix of the ealaxy has been to be estimated2.50&LOeresbem2 and fromthe computed Πα huninosity the (global) star formation rate (SFR) las been derived. which is SFR=O.71\Mow -.," The $\alpha$ flux of the galaxy has been estimated to be $\rm{2.50 \times 53410^{-12}\,erg\,s^{-1}\,cm^{-2}}$ and fromthe computed $\alpha$ luminosity the (global) star formation rate (SFR) has been derived, which is $\rm{SFR = 5350.71\,M_{\sun}\,yr^{-1}}$ ."536" This southern edgeou spiral is slightly larger than the EFOSC? field of view which is BisSHES,", This southern edge–on spiral is slightly larger than the EFOSC2 field of view which is $5\farcm8 \times 5\farcm8$.537 22531 is secu perfectly edge.on., 2531 is seen perfectly edge–on.538" Ta our Πα image alinost no extraplanar tse cutission has beeu detected,", In our $\alpha$ image almost no extraplanar diffuse emission has been detected.539 One fbuneut. (the ChimneyTike feature) is clearly seen. emergiug from the disk radius at A kkpe south of the plane into the halo 6866 Fig.," One filament (the chimney–like feature) is clearly seen, emerging from the disk radius at $R$ kpc south of the plane into the halo (see Fig."540 3 ).," \ref{F3}541 )."542 This feature is marked in Fig., This feature is marked in Fig.543" 3) with a circle,", \ref{F3} with a circle.544 It reaches a height of :—2Xkkpe above the galactic plane., It reaches a height of $z$ kpc above the galactic plane.545 The Πα nage looks pretty much like a strine of pearls., The $\alpha$ image looks pretty much like a string of pearls.546 Several disk regionscan be ideutified. but only the largest ave surrounded by DIG. which is probably not," Several disk regionscan be identified, but only the largest are surrounded by DIG, which is probably not"54710 “bootstrap” [luxes of 2005|403 at these frequencies to rose given in the University of Michigan Radio Astronomy Observatory (UMIUXO) flux database shows that. where closely contemporancous Observations are. available. the agreement with UMICAO Iluxes at these frequencies is within La.,"the “bootstrap” fluxes of 2005+403 at these frequencies to those given in the University of Michigan Radio Astronomy Observatory (UMRAO) flux database shows that, where closely contemporaneous observations are available, the agreement with UMRAO fluxes at these frequencies is within $1\%$."548 Given the excellent agreement at these two frequencies. it is reasonable to assume that the Dux scales at the other observed frequencies are good.," Given the excellent agreement at these two frequencies, it is reasonable to assume that the flux scales at the other observed frequencies are good."549 Aside from the initial amplitude calibration of the ALERLIN observations. the calibration ancl subsequent imaging of the data were carried out using the NILAO software package.," Aside from the initial amplitude calibration of the MERLIN observations, the calibration and subsequent imaging of the data were carried out using the NRAO software package."550 The complex antennae gains were initially clerived for 2005|403 and then interpolated (phase-referenced) to the observations of WIULI146., The complex antennae gains were initially derived for 2005+403 and then interpolated (phase-referenced) to the observations of 146.551 Lo addition. several iterations of phase-only self-calibration were used to refine the antennae gains during the ALERLIN observations of W1t1146 to improve the dynamic range of the final synthesized images.," In addition, several iterations of phase-only self-calibration were used to refine the antennae gains during the MERLIN observations of 146 to improve the dynamic range of the final synthesized images."552 At 1.5 and 4.9 Cillz the VLA data reveal an unresolved source. whereas at 8.4 (11 the radio emission is marginallv resolved.," At 1.5 and 4.9 GHz the VLA data reveal an unresolved source, whereas at 8.4 GHz the radio emission is marginally resolved."553 Llowever. WIULIId6 is well resolved. into a double radio source at 22 CGlLIz.," However, 146 is well resolved into a double radio source at 22 GHz."554 The final synthesized 22-Cillz image is shown in Fig. 2.., The final synthesized 22-GHz image is shown in Fig. \ref{fig:22GHz}.555 We identify the northern and southern components as No» and So»., We identify the northern and southern components as $_{22}$ and $_{22}$.556 The visibility cata suggest both these components are resolved. (Figure 3))., The visibility data suggest both these components are resolved (Figure \ref{fig:22ghz_visibs}) ).557 Gaussian model fits to the visibilities give the diameter of the radio, Gaussian model fits to the visibilities give the diameter of the radio558Note that &* is also a function of r and Equation (31)) only applies when a«1.,Note that $k_z^*$ is also a function of $r$ and Equation \ref{eq:kzcrit}) ) only applies when $\alpha < 1$.559 The absolute stability of the continuum is guaranteed for k- larger than the maximum value of A> in the equilibrium., The absolute stability of the continuum is guaranteed for $k_z$ larger than the maximum value of $k_z^*$ in the equilibrium.560 Figure 2. shows the thermal continuum computed from Equation (28)) for our equilibrium, Figure \ref{fig:continuum} shows the thermal continuum computed from Equation \ref{eq:cont}) ) for our equilibrium561The three high-density regious in the cluster mass distribution. as traced by the hot XN-rav gas (seo Fie. 53).,"The three high-density regions in the cluster mass distribution, as traced by the hot X-ray gas (see Fig. \ref{x_3sig}) ),"562 ave populated by passive galaxies exclusively., are populated by passive galaxies exclusively.563 All the star-forming galaxies noticcably avoid these regions., All the star-forming galaxies noticeably avoid these regions.564 A few ealaxies with narrow cuuission lines are observed close to the southern edge of the southern clump (see Fig. 8))., A few galaxies with narrow emission lines are observed close to the southern edge of the southern clump (see Fig. \ref{el_membs}) ).565" These galaxies are about ((~ 230 kpe) from the centroid of the southeru N-rav substructure, which corresponds to a region where the local ICA density is about one-third of the central density of the southern chump. as derived from the Nay surface brightuess profile."," These galaxies are about $\sim$ 230 kpc) from the centroid of the southern X-ray substructure, which corresponds to a region where the local ICM density is about one-third of the central density of the southern clump, as derived from the X-ray surface brightness profile."566 Iu ters of spectral classification. passive galaxies are split iuto two main groups: k type galaxies (see Dressler et al.," In terms of spectral classification, passive galaxies are split into two main groups: k type galaxies (see Dressler et al."567 1999: Pogeiauti ct al., 1999; Poggianti et al.568" 1999) aud post-starburst ealaxies (Dressler Coma 1983). also called ""E|A” ealaxies."," 1999) and post-starburst galaxies (Dressler Gunn 1983), also called “E+A” galaxies."569 Based on the streneth of the IL;((A £102). line.," Based on the strength of the $_{\delta}$ $\lambda$ 4102) line,"570"quiet region considered where emission is nearly absent has the value of about 4100 erg s! cm"" ?sr-!.",quiet region considered where emission is nearly absent has the value of about 4100 erg $^{-1}$ $^{-2}$ $^{-1}$.571" The SERTS-97 observation at pointing 1 showed variations of intensities in “quiet region"" from 4500 to 12000 erg s-lem-?sr-! over the -long slit."," The SERTS-97 observation at pointing 1 showed variations of intensities in “quiet region"" from 4500 to 12000 erg $^{-1}$ $^{-2}$ $^{-1}$ over the $^{''}$ -long slit."572 DelZanna&Andretta showed that the irradiance measured during (2011)this solar minimum (2006-2008) is about a factor of 1.3 smaller than that measured in 1998 at the beginning of the last solar cycle., \citet{del11} showed that the irradiance measured during this solar minimum (2006–2008) is about a factor of 1.3 smaller than that measured in 1998 at the beginning of the last solar cycle.573 The dispersion of intensities shown in Table 5 suggests that one should be cautious about using the quiet-Sun line as a standard light source., The dispersion of intensities shown in Table \ref{tabhe} suggests that one should be cautious about using the quiet-Sun line as a standard light source.574" In addition, we may verify the EUNIS-07 radiometric calibration by comparing the measured quiet-Sun 304 iintensity with the radiance converted from irradiance measurements (the full-disk flux measured at Earth from this line or a narrow waveband centered at 304 during the solarminimum."," In addition, we may verify the EUNIS-07 radiometric calibration by comparing the measured quiet-Sun 304 intensity with the radiance converted from irradiance measurements (the full-disk flux measured at Earth from this line or a narrow waveband centered at 304 ) during the solarminimum."575" Since the lines have À))negligible limb-brightening and off-limb contribution (DelZanna&Andretta 2011), the conversion for the quiet Sun can be simplymade by the relation (Warrenetal.1998) where is the irradiance, Ro is the solar radius, R is the Fj,Earth-Sun distance, and Ig, is the intensity at disk-center."," Since the lines have negligible limb-brightening and off-limb contribution \citep{del11}, the conversion for the quiet Sun can be simplymade by the relation \citep{war98}576 where $F_{qs}$ is the irradiance, $R_{\sun}$ is the solar radius, $R$ is the Earth-Sun distance, and $I_{qs}$ is the intensity at disk-center."577" When the irradiance uses the unit of photons s! cm-? and the intensity is in erg sσι ια, the conversion coefficient between Fy, and I4, at 304 iis 1.04x10° sr photons erg-!."," When the irradiance uses the unit of photons $^{-1}$ $^{-2}$ and the intensity is in erg $^{-1}$ $^{-2}$ $^{-1}$, the conversion coefficient between $F_{qs}$ and $I_{qs}$ at 304 is $\times$ $^{6}$ sr photons $^{-1}$ ."578" From the quiet-Sun intensity of 4960 erg s-!cm-?sr! by EUNIS-07, it derives irradiance to be 52x10? photons s! cm?, which"," From the quiet-Sun intensity of 4960 erg $^{-1}$ $^{-2}$ $^{-1}$ by EUNIS-07, it derives irradiance to be $52\times10^8$ photons $^{-1}$ $^{-2}$ , which"579work.,work.580 Their results are consistent with the ones presented here., Their results are consistent with the ones presented here.581 Analvziug the secondary eclipse data. we report here the detection of a secondary eclipse aud draw conclusions about the thermal emüssion of LI36bb. aud. refine its orbital parameters. allowing a better πιοταιπας of L136 dynamics by exploring the coutingency of a supplementary plauct.," Analyzing the secondary eclipse data, we report here the detection of a secondary eclipse and draw conclusions about the thermal emission of b and refine its orbital parameters, allowing a better understanding of 436 dynamics by exploring the contingency of a supplementary planet."582 Iu addition. we report here ou additional ground based observations to determine the stellar rotational period.," In addition, we report here on additional ground based observations to determine the stellar rotational period."583 We followed the photometric intensity and the Ca IT IT|KR activity iudex of 1036., We followed the photometric intensity and the Ca II H+K activity index of 436.584 Although the photometric data are sparse aud cover only 50 davs. we find some evidence that the stellar rotational period is of the order of 50 days. which is also cousisteut with loug-terui Call micasurcmicuts.," Although the photometric data are sparse and cover only 50 days, we find some evidence that the stellar rotational period is of the order of 50 days, which is also consistent with long-term CaII measurements."585 Section 2 describes the observations and the reduction procedure., Section 2 describes the observations and the reduction procedure.586 Our analysis of the obtained secondary eclipse time series is described in Section 3., Our analysis of the obtained secondary eclipse time series is described in Section 3.587 Iu Section [. we analyze the infrared cussion from the planet and draw sole couclisions about its atmosphere composition.," In Section 4, we analyze the infrared emission from the planet and draw some conclusions about its atmosphere composition."588 We detail an orbital analysis. enconipassimg the possibility of a perturbing planet. stellar activity aud bb orbital paramcters refinements in Section 5.," We detail an orbital analysis, encompassing the possibility of a perturbing planet, stellar activity and b orbital parameters refinements in Section 5."589 Our couchisions are presented in Section 6., Our conclusions are presented in Section 6.590 L136 has been observed on June 30th UT for 6 hours. ο cover the secondary eclipse. resulting m 19920 frames.," 436 has been observed on June 30th UT for 6 hours, to cover the secondary eclipse, resulting in 49920 frames."591 Observations were made so as to enconipass the expected secondary eclipse window. whose tinue caleulatious were uade by taking iuto account transit timing and orbital eccentricity.," Observations were made so as to encompass the expected secondary eclipse window, whose timing calculations were made by taking into account transit timing and orbital eccentricity."592 Due to the uncertainties on ecceutricitv iud argunenut of poeriastron. a larger time-window was chosen o eusure the detection of the secondary eclipse.," Due to the uncertainties on eccentricity and argument of periastron, a larger time-window was chosen to ensure the detection of the secondary eclipse."593 Data acquisition was made using TRAC in its 58-442 band with he same mode and strategy enmiploved for the primary ransit (CCUOTb)., Data acquisition was made using IRAC in its $\mu$ m band with the same mode and strategy employed for the primary transit (G07b).594 We combine cach set of 61 images using a Dg clipping to get rid off transicut eveuts m the pixel exid. vielding τοῦ stacked nuages for the secondary eclipse. with a temporal sampling of ~ 28s.," We combine each set of 64 images using a $\sigma$ clipping to get rid off transient events in the pixel grid, yielding 780 stacked images for the secondary eclipse, with a temporal sampling of $\sim$ 28s."595 Helioceutric Julian Dax CIID) couversion was made according to the mean Spitzer orbital position at the time of cach exposure aud L136 apparent position., Heliocentric Julian Day (HJD) conversion was made according to the mean $Spitzer$ orbital position at the time of each exposure and 436 apparent position.596 Spitrer position cplemerides were obtained through JPL-Iorizous web interface (9) aud converted frou TT (Terrestrial Dynamic Time) to UTC., $Spitzer$ position ephemerides were obtained through JPL-Horizons web interface \citep{Giorgini:1996ai} and converted from TT (Terrestrial Dynamic Time) to UTC.597 We faced the same instrumental rise issue noticed iu our work on primary transit., We faced the same instrumental rise issue noticed in our work on primary transit.598 To mitigate its effect. we zero weight the eclipse and the first LOO points of the time-serics.," To mitigate its effect, we zero weight the eclipse and the first 100 points of the time-series."599 We then divide the liehteurve bv the best fitting asvinptotic function with three free parameters aud evaluate the average flux outside the eclipse to normalize the ine series. exactly as for the primary transit.," We then divide the lightcurve by the best fitting asymptotic function with three free parameters and evaluate the average flux outside the eclipse to normalize the time series, exactly as for the primary transit."600" The of the resulting time series evaluated outside the eclipse is the same as for the primary (COT): 0.7 παπα, which is 1.2 times 1136s photon noise."," The of the resulting time series evaluated outside the eclipse is the same as for the primary (G07b): 0.7 mmag, which is 1.2 times 436's photon noise."601 To assess the variability of the star. we observed L136 with the Euler Swiss telescope located at La Silla Observatory (Chile) aud the Fraucoois-Navicr Bagnoud Observatorv's (OFXB) 0.6120. telescope located at Saint-Luc (Switzerland).," To assess the variability of the star, we observed 436 with the Euler Swiss telescope located at La Silla Observatory (Chile) and the Françoois-Xavier Bagnoud Observatory's (OFXB) 0.6m telescope located at Saint-Luc (Switzerland)."602 Observations occurred in Ll nights from Mav Ith to May. 21th., Observations occurred in 14 nights from May 4th to May 21th.603 A sequence of LO exposures was done every night., A sequence of 10 exposures was done every night.604 The same strategy used for our 6servation ofthe May 2ud transit (0Τα) was applied (V-baud filter. SOs exposure time. defocus to ~ 97).," The same strategy used for our observation of the May 2nd transit (G07a) was applied (V-band filter, 80s exposure time, defocus to $\sim$ 9”)."605 The data reduction was also simular., The data reduction was also similar.606 We also use for our analysis of the 1136. variability the Mav δις out-of-trausit data and the photometric lightcurves obtained with the OFNB 0.6120 telescope curing our search for the transits of bb (CO7a)., We also use for our analysis of the 436 variability the May 2nd out-of-transit data and the photometric lightcurves obtained with the OFXB 0.6m telescope during our search for the transits of b (G07a).607 We scale OEXD points with Euler ones because of the filters slightly differeut baudpasses., We scale OFXB points with Euler ones because of the filters slightly different bandpasses.608 At the eud. our data amounts to 21 poiuts spamming Ls dave.," At the end, our data amounts to 24 points spanning 48 days."609 The Leltcurve is represented in Fie., The lightcurve is represented in Fig.610 6. aud. discussed iu Sect.," 6, and discussed in Sect."611 5.3., 5.3.612 Since the discovery of hb) (7). we obtained additional spectra of the star with the ESO spectrograph (?3..," Since the discovery of b \citep{Butler:2004dq}, , we obtained additional spectra of the star with the ESO spectrograph \citep{Mayor:2003pb}."613 is mounted on ESO 3.611 telescope aud is dedicated to Ligh precision racial-velocity measurements thanks to its resolution of 110000 and a wavelength range coverage between 3800 aud6800À., is mounted on ESO 3.6m telescope and is dedicated to high precision radial-velocity measurements thanks to its resolution of 110'000 and a wavelength range coverage between 3800 and.614. To assess the stellar activity and rotation we used 23 high SNR spectra from which we measured the Ca IW Iv index., To assess the stellar activity and rotation we used 23 high SNR spectra from which we measured the Ca II H+K index.615 Results are discussed in Sect., Results are discussed in Sect.616 5.3 We fit a non-Iub-darkenued eclipse profile to the secondary eclipse data using the ? algorithm., 5.3 We fit a non-limb-darkened eclipse profile to the secondary eclipse data using the \citet{Mandel:2002wd} algorithm.617 The eccentricity of the orbit is considered as described in G07). taking the values for the eccentricity e aud the aremment of periastrou uw from ALO7.," The eccentricity of the orbit is considered as described in G07b, taking the values for the eccentricity $e$ and the argument of periastron $\omega$ from M07."618 The forumla counecting w to the true anomaly f at the orbital location of the secondary eclipse is: We fix the stellar aud orbital parameters to the values mentioned m Cra., The formula connecting $\omega$ to the true anomaly $f$ at the orbital location of the secondary eclipse is: We fix the stellar and orbital parameters to the values mentioned in G07a.619 The free parameters are the ceutra epoch of the secondary eclipse T; aud the fiux decremeut AF..., The free parameters are the central epoch of the secondary eclipse $T_s$ and the flux decrement $\Delta F_s$.620 The fit procedure aud the error lars estimation is sinular to the oue described in GOT)., The fit procedure and the error bars estimation is similar to the one described in G07b.621 The obtaiue« value for Ty aud AF. ποπιο their respective error bars are eiven in Table 1.," The obtained value for $T_s$ and $\Delta F_s$, including their respective error bars are given in Table 1."622 Figure 1 shows the best-fit theoretical curvesuperimposed ou the lighteuve (zoonie«, Figure 1 shows the best-fit theoretical curvesuperimposed on the lightcurve (zoomed623The FIR luminosity is often used as a measure of the current star formation rate (SFR). since 1t is assumed that FIR emission is mainly due to dust heating by massive young stars.,"The FIR luminosity is often used as a measure of the current star formation rate (SFR), since it is assumed that FIR emission is mainly due to dust heating by massive young stars."624" The total IR luminosity of iis 1.3x10"" eeress7!. according to the precepts of Draine&Li(2007). and with fluxes from Daleetal.(2005)."," The total IR luminosity of is $\times10^{44}$ $^{-1}$, according to the precepts of \citet{draine07} and with fluxes from \citet{dale05}."625.. This corresponds to à SFR of -6 MM. yyr! (Kennicutt1998)., This corresponds to a SFR of $\sim6$ $_\odot$ $^{-1}$ \citep{kennicutt98}.626. In the bulge of3627.. there is little observed SF 1994;Reganetal. 2002).. and the SFR given by wwithin a nuclear region of diameter ~ iis 0.078M..yr! (Reganetal.2002)... ~3 times lower than found in the bar itself. and ~4 times lower than the spiral arms (Reganetal.2002).," In the bulge of, there is little observed SF \citep[][]{smith94,regan02}, and the SFR given by within a nuclear region of diameter $\sim$ is $^{-1}$ \citep[][]{regan02}, $\sim$ 3 times lower than found in the bar itself, and $\sim$ 4 times lower than the spiral arms \citep[][]{regan02}."627. Part of this deficit in the nuclear -derived SFR may arise from dust extinction. given that the mean Ay in the inner ((diameter) is ~2mmag (Calzettietal.2007).," Part of this deficit in the nuclear -derived SFR may arise from dust extinction, given that the mean $_V$ in the inner (diameter) is $\sim$ mag \citep{calzetti07}."628. In any case. in3627.. most of the SF is extranuclear. along the bar. particularly where it terminates and the spiral arms emerge (see Figs.," In any case, in, most of the SF is extranuclear, along the bar, particularly where it terminates and the spiral arms emerge (see Figs."629 17 and 18))., \ref{fig:mips70-bima} and \ref{fig:mips160-bima}) ).630 The eemission is confined mainly to the nucleus and the bar. particularly the ansae (see Fig. 17)).," The emission is confined mainly to the nucleus and the bar, particularly the ansae (see Fig. \ref{fig:mips70-bima}) )."631 The eemission (albeit with lower resolution). is more broadly distributed. especially the bar.," The emission (albeit with lower resolution), is more broadly distributed, especially the bar."632 This IR. morphology suggests that the dust along the bar is warmer than around the bar. probably heated by the massive stars in the recent formation episodes.," This IR morphology suggests that the dust along the bar is warmer than around the bar, probably heated by the massive stars in the recent star-formation episodes."633 In galaxies with weak SF activity. (e.g.NGC4736.Smithetal. 1994).. dust heating by non-OB stars may also contribute significantly (e.g..deJongetal.1984;Bothun 1989)..," In galaxies with weak SF activity \citep[e.g. NGC 4736,][]{smith94}, dust heating by non-OB stars may also contribute significantly \citep[e.g.,][]{dejong84,bothun89}."634 This more quiescent heating source may be especially important in the central regions of early/type spiral galaxies with massive bulges and little nuclear or eircumnuclear SF. such as3627.," This more quiescent heating source may be especially important in the central regions of early/type spiral galaxies with massive bulges and little nuclear or circumnuclear SF, such as."635. The ratio of FIR to luminosity for the bulge of us of «8100. significantly larger than for the star-forming regions in this galaxy. between -1000 and -2000.," The ratio of FIR to luminosity for the bulge of is of $\sim$ 8100, significantly larger than for the star-forming regions in this galaxy, between $\sim$ 1000 and $\sim$ 2000."636 The L(FIR)/L(Ha)) ratio is also higher than can be accounted for by obscured SF with a normal initial. mass function. using extinction measurements derived from CO(1-0) and FIR data (Smithetal.1994).," The ) ratio is also higher than can be accounted for by obscured SF with a normal initial mass function, using extinction measurements derived from $^{12}$ CO(1–0) and FIR data \citep[][]{smith94}."637. Thus. the older stars probably contribute significantly to the dust heating in the bulge of citep[]| |smith94..," Thus, the older stars probably contribute significantly to the dust heating in the bulge of \\citep[][]{smith94}."638 A low nuclear SER 1s consistent with the CO/HCN ratio (10) discussed in Sect. 3.., A low nuclear SFR is consistent with the CO/HCN ratio (10) discussed in Sect. \ref{sec:30m}.639" Higher ratios suggest that excitation by SF is dominant over AGN excitation in the cireumnuclear region. but we found a ""normal"" CO/HCN ratio for3627.. not surprisingly given its low SER."," Higher ratios suggest that excitation by SF is dominant over AGN excitation in the circumnuclear region, but we found a “normal” CO/HCN ratio for, not surprisingly given its low SFR."640 The gravitational torques derived from the stellar potential in the inner region of aallow to account for the gas kinematies derived from CO and examine the efficiency of gravitational torques exerted on the gas., The gravitational torques derived from the stellar potential in the inner region of allow to account for the gas kinematics derived from CO and examine the efficiency of gravitational torques exerted on the gas.641 As described in previous NUGA papers (e.g..García-Burilloetal. 2005).. to compute the gravitational torques we assume that NIR images give the best approximation for the total stellar mass distribution. being less affected than optical images by dust extinction or stellar population bias.," As described in previous NUGA papers \citep[e.g.,][]{santi05}, to compute the gravitational torques we assume that NIR images give the best approximation for the total stellar mass distribution, being less affected than optical images by dust extinction or stellar population bias."642 We computed the torques using both HST--NICMOS F160W and Spitzer--IRAC limages., We computed the torques using both -NICMOS F160W and -IRAC images.643 They yield complementary. results. the torques computed from the HST--NICMOS FI60W image compared," They yield complementary results, the torques computed from the -NICMOS F160W image compared"644possible since Chere were no other sources in the field.,possible since there were no other sources in the field.645 We estimate the telescope pointing accuracy to be less than 5 aresec., We estimate the telescope pointing accuracy to be less than 5 arcsec.646 Fig., Fig.647 shows a map of the IH» emission towards ΗνΛο 16547—4247., \ref{plotone} shows a map of the $_2$ emission towards IRAS $-$ 4247.648 There is a complex chain of emission with three major concentrations (labeled A to C)., There is a complex chain of emission with three major concentrations (labeled A to C).649 Several of the brightest emission knots within each concentration have been labeled and their coordinates and fIuxes are given in Table 1.., Several of the brightest emission knots within each concentration have been labeled and their coordinates and fluxes are given in Table \ref{tableone}.650 The projected distance between the two outermost knots (At and C2) is 110 aresee (1.5 pe at the distance of 2.9 kpc. DronBnan private communication).," The projected distance between the two outermost knots (A4 and C2) is 110 arcsec (1.5 pc at the distance of 2.9 kpc, Bronfman private communication)."651 Both of (hese knots are approximately svimmetrically offset [rom the raclio jet detected by (2003).., Both of these knots are approximately symmetrically offset from the radio jet detected by \citet{Garay03}. .652 No emission arising from the Dr5 line was detected., No emission arising from the $\gamma$ line was detected.653 The II» emission has the morphological characteristics of HII objects arising from the interaction of a collimated flow with the ambient medium., The $_2$ emission has the morphological characteristics of HH objects arising from the interaction of a collimated flow with the ambient medium.654 Concentration A has several knots in the shape of bow-shocks. all pointing away from the direction of the radio jet.," Concentration A has several knots in the shape of bow-shocks, all pointing away from the direction of the radio jet."655 Their arrangement is consistent with an elongated outflow cavity., Their arrangement is consistent with an elongated outflow cavity.656 Another series of emission knots may exist. further north but is difficult to distinguish against the artifacts from the bright star., Another series of emission knots may exist further north but is difficult to distinguish against the artifacts from the bright star.657 The morphology of concentration D is more complex and consists of two main enission structures., The morphology of concentration B is more complex and consists of two main emission structures.658 Both structures appear to delineate flows originating from a direction that is skewecl from the location of the radio jet: one in à northeast.southwest direction (Bl. D2. and BG) and one in a north.south direction (D3. Bt and D5).," Both structures appear to delineate flows originating from a direction that is skewed from the location of the radio jet: one in a northeast–southwest direction (B1, B2, and B6) and one in a north–south direction (B3, B4 and B5)."659 This may be evidence of additional oulllows [rom less-massive stars or an indication of precession of the flow originating from the detected radio jet., This may be evidence of additional outflows from less-massive stars or an indication of precession of the flow originating from the detected radio jet.660 There are fewer bright emission knots in concentration C and the morphology is less well-defined., There are fewer bright emission knots in concentration C and the morphology is less well-defined.661 There are a couple of Taint filaments in the shape of pointing away [rom the direction of the radio jet., There are a couple of faint filaments in the shape of bow-shocks pointing away from the direction of the radio jet.662 These are most likely part of the counter-flow to concentration A. Fie., These are most likely part of the counter-flow to concentration A. Fig.663 3 illustrates the comparison between the IH» emission (without any continuum subtraction) and (a) the 1.2-mm dust continuum emission and (b) the 8.6-Gllz continuum enission (taken from Qarayetal.(2003)., \ref{plottwo} illustrates the comparison between the $_2$ emission (without any continuum subtraction) and (a) the 1.2-mm dust continuum emission and (b) the 8.6-GHz continuum emission taken from \citet{Garay03}.664.. The actual thermal radio jet corresponds to the brightest 8.6-GlIIz emission component., The actual thermal radio jet corresponds to the brightest 8.6-GHz emission component.665 There is a fainter source offset to (he southeast and whose spectral index is not known., There is a fainter source offset to the southeast and whose spectral index is not known.666 It is not certain what role (if anv) this source plavs., It is not certain what role (if any) this source plays.667 The chain of Hà emission is oriented in the same direction as (he triple radio source and is contained within the molecular core traced by the 1.2-imm continuum emission., The chain of $_2$ emission is oriented in the same direction as the triple radio source and is contained within the molecular core traced by the 1.2-mm continuum emission.668 One of theII emission knots (Bl) is associated will the southern non-thermal radio component., One of the$_2$ emission knots (B1) is associated with the southern non-thermal radio component.669any given epoch.,any given epoch.670 As an illustrative case. we have shown iu figure 3. the spectral energy distributions corrected for Ay=0.09 of SMC type dust.," As an illustrative case, we have shown in figure \ref{sedfig} the spectral energy distributions corrected for $A_V = 0.09$ of SMC type dust."671 This correction gives the minimum \7/cof (1.15 for I degrees of [reedoim)) for the residuals of the plotted spectral energy. distributious relative to the spectral slopes reported iu table 3.., This correction gives the minimum $\chi^2 / \dof$ $1.45$ for 4 degrees of freedom) for the residuals of the plotted spectral energy distributions relative to the spectral slopes reported in table \ref{slopetable}.672 LAIC extinction does mareinally worse than SAIC extinction. whileH any zunouut ofH MilkyH Way- extinctionH inH the host clegrades V7.," LMC extinction does marginally worse than SMC extinction, while any amount of Milky Way extinction in the host degrades $\chi^2$."673> We- have used the analyticH extinction law fitting forms of Pei (1992) in deriving these estimates., We have used the analytic extinction law fitting forms of Pei (1992) in deriving these estimates.674 Jeuseu et al (2000) have applied a similar analysis iucorporatiug optical spectra as well as broadband colors., Jensen et al (2000) have applied a similar analysis incorporating optical spectra as well as broadband colors.675 They also fiud a siguificautly better fit for SMC extinetion than for eitherMW or LMC extinction. auc derive Ay=O1LL40.01 magnitudes for the SMC model.," They also find a significantly better fit for SMC extinction than for eitherMW or LMC extinction, and derive $A_V = 0.14 \pm 0.01$ magnitudes for the SMC model."676" Usine our multiple epoch SEDs. we find a 2,. ⋅ ⋅ ⋅ ⋅ ⊳∖∩∐↕≺↵∖∖↽∐⋜∐∖∖↽∩↕⋅↜∖≺↵∖−↙∕∕≼⇂⋅∩⋅↥∶⊔∫↖∖↥∩↕⋅↕∐↩∐⋅≺↵⊸∖⋃∐∢∙⊔∩↥≺↵⊳∖∐↕⋜↕↕≺↵↕∐⋜↕∐↥∩↕⋅∩⋃⋅⊳∖⋅∣≻⋯↕⇂≺↵↕∖∖↽∩↕⋅≺↲⊳∖⋃∐⊳∖ ⋜⋃⋅≺↵↥↽∐⋅∩∣≻⋜≹∣≻↥⊽∖⇁∢∙∩∐⊳∖↥⊳∖↕≺↵∐↕∖∖↽∐∐↥∐↕∐↩≺↵⋅⋅∩↓∷∖⋅↩⊳∖↥↽≻≺↵∢∙↥⋜↕∐⊽∖⊽∐∎↕≺↵↕⋅≺↲⋜↕↓⋅≺↵⋯∏≺⇂≺↵∐↕∐∎∐↵≺⇂⊳∖⊽∖⇁⊳∖↕≺↵⋯⋜↕∏∢∙≺↲⊔⋅∩↥⋅⊳∖∩↥∎ ≧⋋↴∪⋅∪↖∖⋝⋯⋜↕∑≟∐∐⋯⇂↩↕∐↕∐≺↵↥↽≻∐"," Using our multiple epoch SEDs, we find a somewhat worse $\chi^2 / \dof = 1.98$ for their extinction estimate than for ours, but the two results are probably consistent within the errors, especially if there are unidentified systematic errors of $\ga 0.08$ magnitude in the photometry (see section \ref{colsec}) )."677∩↕∩⋯≺↵⋃⋅⊽∖⇁↸⋮⊳∖≺↵≺↵↠∖≺↲∢∙∏∩⊔⋅↴⊔⋅ ⊺∐≺↵⋜↕↥↽≻↥↽≻⋜⋃⋅≺↵∐↕∖∖↽≺↵⋜↕↕⊆∐≺↲↜∖↠∖∩↥∎↕∐≺↲⊇↽⋡⊤⋅↴⇀−∖ ∐∎≺↵⋜↕⊓⊔⋅≺↵↥∐↕∐↩↩⊸∖∏∐∢∙∏∩∐∢∙⋯⋅∖⊽≺↵∩⊓∐≺↵∐∩⊳∖↕∑∸⋜↕↥⋜∟∖⊽∖⊽↥⊳∖↕⋅≺↵∐∐∐↥⊳∖∢∙≺↲∐↕∩↥∎≺⇂⋃⊳∖↕⋜↕⊓≺↵⋯⇂⋜↕∏∩∐↥⋜↕∖∖⋱∖↥∎∩↓⋅⋜↕∢∙∏∖⇁≺↵↥⊽∖⇁ ⊳∖↕⋜↕↓⋅↥∎≺∐⋅∐∏∐∑≟∑≟⋜↕⋜∟∖↥↩⊳∖↸⋮≺↵⋅∑≟⋅⋅↕∐↩↥∖↓⋜↕∑≟≺↵∐⋜↕∐↥∢∙∊↽⊲↥∩⋯⇂↠∖∐⊃≺↲↥⊽∐⋯⊇⋜↕∐≺⊔⋅≺↵↥∎≺↵↕⋅≺↲∐∢∙≺↵⊳∖↕∐≺↵↓⋅≺↵↥∐↴⋜↕∐≺⇂⊳∖↕⋜↕↓⋅∣∥⊔⋅⊳∖↕ ealaxies [Gordon. Calzetti. Witt 1997]).," The apparent weakness of the $2175$ feature in the extinction curve of the host galaxy is reminiscent of dust attenuation laws for actively star forming galaxies (e.g., the Magellanic Clouds [Pei 1992 and references therein] and starburst galaxies [Gordon, Calzetti, Witt 1997])."678 This may be further cireumstantial evidence linking GBRBs to actively star forming ealaxies., This may be further circumstantial evidence linking GRBs to actively star forming galaxies.679 Alternatively. such au extinction law mieht be observed if GRBs preferentially destroy the sinall carbonaceous particles thought to be carriers of the 2175 feature. but this explanation would only work if much of the dust optical depth arises near the uiaximnunr radius where the burst cau destroy graius WWasman Draine 2000: Fruchter. νο. Rhoads 2000).," Alternatively, such an extinction law might be observed if GRBs preferentially destroy the small carbonaceous particles thought to be carriers of the $2175$ feature, but this explanation would only work if much of the dust optical depth arises near the maximum radius where the burst can destroy grains Waxman Draine 2000; Fruchter, Krolik, Rhoads 2000)."680 Iu order to determine pliysical parameters of the afterglow. we need to measure the peak fIux density aud the locatious of breaks in the afterglow spectrum (Wijers Calama 1999).," In order to determine physical parameters of the afterglow, we need to measure the peak flux density and the locations of breaks in the afterglow spectrum (Wijers Galama 1999)."681 We now do this (iusolar as possible) by combinine our optical-IR spectral slope measurements with the sttbuatllimeter aud radio data., We now do this (insofar as possible) by combining our optical-IR spectral slope measurements with the submillimeter and radio data.682" We are lookingfor four uumbers: The frequency riposo, aud fux deusity fjiax at the peak in f/,: the cooling frequency (. aud the self-absorptiou [requeney 7abs:"," We are lookingfor four numbers: The frequency $\nu_{hbox{max}}$ and flux density $f_{\nu,\hbox{max}}$ at the peak in $f_\nu$; the cooling frequency $\nu_c$ , and the self-absorption frequency $\nu_{\hbox{abs}}$."683 The spectral slope is expected to be —p/2 for v>νο. —(p—1)/2 for max€rb<re +1/3 for MapsVEXVinax- and +2 for v<Vols (Sari. Piran. Narayan 1998).," The spectral slope is expected to be $-p/2$ for $\nu > \nu_c$, $-(p-1)/2$ for $\nu_{\hbox{max}} < \nu < \nu_c$, $+1/3$ for $\nu_{\hbox{abs}} < \nu <684\nu_{\hbox{max}}$, and $+2$ for $\nu < \nu_{\hbox{abs}}$ (Sari, Piran, Narayan 1998)."685 Here p is the power law iudex of electrons recently accelerated at the external shock of the expanding GRB remnant., Here $p$ is the power law index of electrons recently accelerated at the external shock of the expanding GRB remnant.686 Extrapolating the optical-IR spectral slopes to lower frequencies. we see that a strong spectral break is required wear or above the 250GHz measurement by Bertoldi (2000) on March. L385.," Extrapolating the optical-IR spectral slopes to lower frequencies, we see that a strong spectral break is required near or above the $250687\GHz$ measurement by Bertoldi (2000) on March 4.385."688" The radio data from March 5.67 (Berger et al 2000) are compatible with fj,xvt between 22 aud 250GHz. so we determine rax aud [μας by extrapolating this beliavior until it intersects the extrapolation [rom optical-IR. data (see figure £))."," The radio data from March 5.67 (Berger et al 2000) are compatible with $f_\nu689\propto \nu^{1/3}$ between $22$ and $250 \GHz$, so we determine $\nu_{\hbox{max}}$ and $f_{\nu,\hbox{max}}$ by extrapolating this behavior until it intersects the extrapolation from optical-IR data (see figure \ref{sedwide}) )."690 Using fluxes corrected for both Galactic dust and Ay=0.009 magnitude of SMC type extinction at 2=) 2.03. we obtain log(Áax/ and log(iinax/qmdJy) 0.100.050. where the error bars account onlyfor photometric errors ou the data.," Using fluxes corrected for both Galactic dust and $A_V = 0.09$ magnitude of SMC type extinction at $z=2.03$ , we obtain $\log(\nu_{\hbox{max}}/\Hz) = 11.81 \pm 0.10$ and $\log(f_{\nu,\hbox{max}} / \mJy ) = 0.46 \pm 0.05$ , where the error bars account onlyfor photometric errors on the data."691 If we do not apply any correction lor host galaxy extiuction. we instead obtain," If we do not apply any correction for host galaxy extinction, we instead obtain"692objects from PLJ have the same average color is.,objects from PLJ have the same average color is.693 The Nohuogorov-Siniznoy two-sample test applied to the same samples indicates probability that the distributions of 7 ICKBOs and the 31 objects from PLJ have the same average color ls«, The Kolmogorov-Smirnov two-sample test applied to the same samples indicates probability that the distributions of 7 ICKBOs and the 34 objects from PLJ have the same average color is.694 The decrease in sample size of the « ssuuples presunablv accounts for the slight increase m probability for the Μαμά] samples., The decrease in sample size of the $<$ samples presumably accounts for the slight increase in probability for the smaller samples.695 When we began this particular project. we thought that the low aud immer bolt objects were a natural sunward extension of the cold classical population found between 12 <a « IN AU at low candz.," When we began this particular project, we thought that the low and inner belt objects were a natural sunward extension of the cold classical population found between 42 $<$ a $<$ 48 AU at low and."696 Our original aim was simply to measure colors for a sample of these objects to compare the colors with the cold main belt objects., Our original aim was simply to measure colors for a sample of these objects to compare the colors with the cold main belt objects.697 As discussed in Section IL. the dynamical status of the lowc. ¢inner disk objects is a major question.," As discussed in Section II, the dynamical status of the low, inner disk objects is a major question."698 The importance of these objects is ouly now being recognized: for example. the Nice model of Levisonetal.(2008). sid nothing specifically about the inner classical disk.," The importance of these objects is only now being recognized; for example, the Nice model of \citet{LevisonNice08} said nothing specifically about the inner classical disk."699 Mich as we would like to lave a definite answer to the precise relationship between the iuncr-classical IKBOs aud the cold-classical KBOs. neither our color data nor the dvuauical models can resolve this issue at present.," Much as we would like to have a definite answer to the precise relationship between the inner-classical KBOs and the cold-classical KBOs, neither our color data nor the dynamical models can resolve this issue at present."700 One of the ain issues is the dynanmdüeal origin of the ΠΙΟ classical objects. as discussed in section IT above.," One of the main issues is the dynamical origin of the inner classical objects, as discussed in section II above."701" So lone as this issue remains unsettled. we are left with admittedly frustrating result of wo different possible ""conclusions for how our colors nav relate these objects to the other »pulatious."," So long as this issue remains unsettled, we are left with admittedly frustrating result of two different possible “conclusions"" for how our colors may relate these objects to the other populations."702 Though of course more color data are always welcome. and could perhaps provide a stronger statement about the connection. or ack of connection. with the plutiuo population. hat will have to wait for further discoveries of brighter iuner classical objects amenable to ueasurue colors.," Though of course more color data are always welcome, and could perhaps provide a stronger statement about the connection, or lack of connection, with the plutino population, that will have to wait for further discoveries of brighter inner classical objects amenable to measuring colors."703 Iu the imeamvhile. we await urther developments from the dvuaiiicists to xopose a possible resolution to this dileuuna.," In the meanwhile, we await further developments from the dynamicists to propose a possible resolution to this dilemma."704 The R color distribution of a sample of LO iuner disk Euiper Belt objects is shown to be not incousisteut with the color distributiou of salples of plutinos. Centaurs aud scattered disk objects.," The $-$ R color distribution of a sample of 10 inner disk Kuiper Belt objects is shown to be not inconsistent with the color distribution of samples of plutinos, Centaurs and scattered disk objects."705 The current inner disk guple has both red and neutral colored objects. roenuünisceut of the Ceutaurs. which have a bimodal BR color distribution.," The current inner disk sample has both red and neutral colored objects, reminiscent of the Centaurs, which have a bimodal $-$ R color distribution."706 However. we cannot claim that the ICKDOs have a statistically significant biiiodal sienal with the present data.," However, we cannot claim that the ICKBOs have a statistically significant bimodal signal with the present data."707 The average BOR color of the ICKBOs are inconsistent. at the level ov higher. depending on specific sample and statistical test used. with the colors of a sample of cold (low inclination) classical KBOs with sciimajor axes between 12 Hid [8 AU.," The average $-$ R color of the ICKBOs are inconsistent, at the level or higher, depending on specific sample and statistical test used, with the colors of a sample of cold (low inclination) classical KBOs with semimajor axes between 42 and 48 AU."708 Possible couclusious of this are: 1) The inner disk objects we observed. even though they are of low and£.. are members of a population analogous to the classical KBOs.," Possible conclusions of this are: 1) The inner disk objects we observed, even though they are of low and, are members of a population analogous to the classical KBOs."709 The hot classical KBOs do not show a predominately red optical color distribution. as do the cold classical objects.," The hot classical KBOs do not show a predominately red optical color distribution, as do the cold classical objects."710 2) If the iuncr disk objects axe in fact ΠΟΙΟΥΣ of a cold inner disk population which is a suu continuation of the cold classical objects between 12 and [8 AU. then there is a radial color eracdieut iu the colors of this cold disk population of KBOs.," 2) If the inner disk objects are in fact members of a cold inner disk population which is a sunward continuation of the cold classical objects between 42 and 48 AU, then there is a radial color gradient in the colors of this cold disk population of KBOs."711 We thauk the NASA Planetary ΑποώμώιἩ program for financial support of this research aud the NASA Ixeck and Vatican Observatory telescope allocation comunittees for consisteut allocation of telescope time., We thank the NASA Planetary Astronomy program for financial support of this research and the NASA Keck and Vatican Observatory telescope allocation committees for consistent allocation of telescope time.712 We thauk Brett Gladman and J.J. Navelaars for discussions of the immer classical belt which helped motivate this work., We thank Brett Gladman and J.J. Kavelaars for discussions of the inner classical belt which helped motivate this work.713jets currently exist lor specific tests against the data.,jets currently exist for specific tests against the data.714 Additionally. we do not believe that the observational evidence cited in favor of shear is unambiguous: 1) while VLDP maps of the QSO 10554018 and of the DL Lac-tvpe object 08204225 show longitudinal magnetic fields in the parsec-scale flows. we do not believe that there is conclusive evidence (hat these are associated with shear. nor do these two cases explain the QSO0-DL Lac dichotomy since one of the objects is a BL Lac: 2) recent analvsis of VLDP maps for 12 blazars (omanetal.2002) observed during a one-year lime period showed that the EVPAs of the jet features rotated with time such that the magnetic field has become more aligned with the jet axis with increasing distance from the core. but the origin of the change in net field orientation could not be unambiguously identified: it is neither clear whether an increased importance of shear will core distance could produce such changes within the linear dimensions probed. nor was the expected associated increased ordering of the field identified in 4 of the 5 cases.," Additionally, we do not believe that the observational evidence cited in favor of shear is unambiguous: 1) while VLBP maps of the QSO 1055+018 and of the BL Lac-type object 0820+225 show longitudinal sheath-like magnetic fields in the parsec-scale flows, we do not believe that there is conclusive evidence that these are associated with shear, nor do these two cases explain the QSO-BL Lac dichotomy since one of the objects is a BL Lac; 2) recent analysis of VLBP maps for 12 blazars \citep{hom02} observed during a one-year time period showed that the EVPAs of the jet features rotated with time such that the magnetic field has become more aligned with the jet axis with increasing distance from the core, but the origin of the change in net field orientation could not be unambiguously identified: it is neither clear whether an increased importance of shear with core distance could produce such changes within the linear dimensions probed, nor was the expected associated increased ordering of the field identified in 4 of the 5 cases."715 Temporal changes in EVPA are currently being studied within the framework of evolving oblique shock models using VLBA data obtained over a 30-month lime span. allowing a larger range of core distances to be explored. and including a wider range of frequencies (Aller.Aller.&Inehes2002).," Temporal changes in EVPA are currently being studied within the framework of evolving oblique shock models using VLBA data obtained over a 30-month time span, allowing a larger range of core distances to be explored, and including a wider range of frequencies \citep{all02}."716. Precession has been discussed widely in (he literature. aud might be expected based on the range of morphologies seen in maps rom VLD surveys (INellernmannetal.1998).. but we are unable to identilv its expected signature in our data.," Precession has been discussed widely in the literature, and might be expected based on the range of morphologies seen in maps from VLB surveys \citep{kel98}, but we are unable to identify its expected signature in our data."717 A combination of longterm monitoring data and VLBI mapping may lead to more conclusive evidence. and (he data included in a recent study of the BL Lac object ON 231 (Alassaroetal.2001). is an observational initial step in this direction.," A combination of longterm monitoring data and VLBI mapping may lead to more conclusive evidence, and the data included in a recent study of the BL Lac object ON 231 \citep{mas01} is an observational initial step in this direction."718 Hlowever. while evidence supporling (he helical character of jet flows is strong. this may result from instabilities in the flow rather than from precession (ILardee.IIughes.Rosen.andGomez2001).," However, while evidence supporting the helical character of jet flows is strong, this may result from instabilities in the flow rather than from precession \citep{har01}."719. Thus. in our view. a convincing case for precession has only been made [or a single source: OJ 287.," Thus, in our view, a convincing case for precession has only been made for a single source: OJ 287."720 The millimeter imaging data now available. combined with sophisticated analvsis tools. perimit studies of the evolution of the structure in the inner regions of the jet.," The millimeter imaging data now available, combined with sophisticated analysis tools, permit studies of the evolution of the structure in the inner regions of the jet."721 With the aid of such modeling. it is now possible to follow the structural development of individual regions in both polarization and total flux. aud (ο identfv the complex changes in the flow direction. including bends and changes in the underlviug magnetic field direction on scales.," With the aid of such modeling, it is now possible to follow the structural development of individual regions in both polarization and total flux, and to identify the complex changes in the flow direction, including bends and changes in the underlying magnetic field direction on sub-milliarcsecond scales."722 These data probe quite different regions from those initially studied ab 5 Gllz. and both opacity effects and Faraday. effects are reduced.," These data probe quite different regions from those initially studied at 5 GHz, and both opacity effects and Faraday effects are reduced."723 Detailed model fitting in the range 15 to 43 GIIz has already. demonstrated that the so-called ‘core’. may. in fact. contain contributions from newlv-emerging. blended components. further complicating an earlier simpler picture.," Detailed model fitting in the range 15 to 43 GHz has already demonstrated that the so-called `core', may, in fact, contain contributions from newly-emerging, blended components, further complicating an earlier simpler picture."724 Indeed. recent model fitting of the VLBP data for some blazars. of both the BL Lac and QSO classes. demonstrate quantitatively (hat the major contribution to the integrated. polarized flix comes [rom the ‘core’ itself or the," Indeed, recent model fitting of the VLBP data for some core-dominated blazars, of both the BL Lac and QSO classes, demonstrate quantitatively that the major contribution to the integrated polarized flux comes from the `core' itself or the"725an icrease in water abundance.,an increase in water abundance.726 Thus. we focus our attention to the molecular clouds with temperature less than 200Ix.," Thus, we focus our attention to the molecular clouds with temperature less than $200~\mathrm{K}$."727 For this temperature range. the values of parameters Αρ and Jy) are given in the Fig. 1..," For this temperature range, the values of parameters $\Lambda_{(n)}$ and $\beta_{(n)}$ are given in the Fig. \ref{fitlambda}."728 As we know. the thermal process may be important if the dynamical time-scale is in excess of (he cooling time-scale.," As we know, the thermal process may be important if the dynamical time-scale is in excess of the cooling time-scale."729 We consider the contraction time-scale of a evlindrical molecular cloud as a multiple >1 of the free-fall time-scale. where (he latter is written for the spherical uniform density distribution.," We consider the contraction time-scale of a cylindrical molecular cloud as a multiple $\eta \geq 1$ of the free-fall time-scale, where the latter is written for the spherical uniform density distribution."730 The cooling time-scale /=3AT/2im and the Iree-Iall time-scale (2)) areshown in Fig. 2.., The cooling time-scale $t_c \equiv 3kT/2m\Lambda$ and the free-fall time-scale \ref{freetime}) ) areshown in Fig. \ref{timescale}.731 In. [ast contraction (4)~ 1). This important in small densities (e.g. less than 10H.? for T—LOOK). while in slow contraction (7>> 1) in which the cooling time-scale is much smaller than (he contraction Gime-scale. (he importance of TI as a trigger mechanism is niuch evident.," In fast contraction $\eta \sim 1$ ), TI is important in small densities (e.g., less than $10^{14} m^{-3}$ for $T=100\mathrm{K}$ ), while in slow contraction $\eta >>1$ ) in which the cooling time-scale is much smaller than the contraction time-scale, the importance of TI as a trigger mechanism is much evident."732 There are several different heating mechanisms in models of interstellar matters., There are several different heating mechanisms in models of interstellar matters.733 Since the ultraviolet photons are mostly screened out in dense molecular clouds. heating by collisional de-excitation of Πο molecules after radiative excitation. of Lyman bands. photoemission from grains. radiative dissociation of ο. and by chemical reactions is not important.," Since the ultraviolet photons are mostly screened out in dense molecular clouds, heating by collisional de-excitation of $H_2$ molecules after radiative excitation of Lyman bands, photoemission from grains, radiative dissociation of $H_2$, and by chemical reactions is not important."734 In addition. because of the small neutral hydrogen abundance. heating due to ejection of newly lormed // molecules from grain surfaces is negligible.," In addition, because of the small neutral hydrogen abundance, heating due to ejection of newly formed $H_2$ molecules from grain surfaces is negligible."735 The heating due to cosmic ravs with sufficient. energies (~ 100MeV) to penetrate dense clouds is commonly about Tep=25x10“JketsἘν with assumption of an ionization rate per //» molecule of 2x10Ms Land amean energy gains per ionization of 19eV (e.g.. Glassgold and Langer 1973).," The heating due to cosmic rays with sufficient energies $\sim 100736\mathrm{MeV}$ ) to penetrate dense clouds is commonly about $\Gamma_{CR} = 2.5 \times 10^{-8} \mathrm{J.kg^{-1}.s^{-1}}$, with assumption of an ionization rate per $H_2$ molecule of $2\times73710^{-17} \mathrm{s^{-1}}$ and a mean energy gains per ionization of $19 \mathrm{eV}$ (e.g., Glassgold and Langer 1973)."738 Following Black (1987) the turbulence dissipation heating rate can be estimated as where ty) is the turbulent. velocity and / is the eddy seale., Following Black (1987) the turbulence dissipation heating rate can be estimated as where $v_{turb}$ is the turbulent velocity and $l$ is the eddy scale.739 With Όρων~1kins! and {ο1pe. we obtain Fogc1.6x10“keFs !. comparable to half of the heating rate of cosmic ravs.," With $v_{turb} \sim 1~ \mathrm{km.s}^{-1}$ and $l\sim7401~\mathrm{pc}$, we obtain $\Gamma_{TR} \sim 1.6 \times 10^{-8}741\mathrm{J.kg}^{-1}.\mathrm{s}^{-1}$ , comparable to half of the heating rate of cosmic rays."742 In this way. we collect the values of these two heating rates to obtain ~Ll1x10kgts !.," In this way, we collect the values of these two heating rates to obtain $\sim 4.1\times 10^{-8}743\mathrm{J.kg}^{-1}.\mathrm{s}^{-1}$ ."744 Another important heating mechanism of sell-gravitating contracting cloud is the heating produced by gravitational compression work., Another important heating mechanism of self-gravitating contracting cloud is the heating produced by gravitational compression work.745 Anestimation lor this, Anestimation for this746blending is the line at 74.860 wwhich contains Mg VIII and Fe XIII.,blending is the line at 74.860 which contains Mg VIII and Fe XIII.747 We have measured line ratios of density-sensitive He-like triplets from the LETGS and RGS spectra. taking into account the photo-exciting UV flux (Porquet et al.," We have measured line ratios of density-sensitive He-like triplets from the LETGS and RGS spectra, taking into account the photo-exciting UV flux (Porquet et al."748 2001)., 2001).749 Our results are consistent in both instruments (ων.~1029 7) and similar to those of Ness et al. (, Our results are consistent in both instruments $n_\mathrm{e} \approx 10^{10}$ $^{-3}$ ) and similar to those of Ness et al. (7502001) and our values given in Table 4.,2001) and our values given in Table 4.751 These results are also comparable to values obtained by Schrijver et al. (, These results are also comparable to values obtained by Schrijver et al. (7521995) and Schmitt et al. (,1995) and Schmitt et al. (7531996) and to values for the Sun (Drake et al.,1996) and to values for the Sun (Drake et al.754 The RGS and LETGS spectra of the corona of Procyon below 40 aare dominated by the H- and He-like transitions of C. N. and O and by Fe XVII lines.," The RGS and LETGS spectra of the corona of Procyon below 40 are dominated by the H- and He-like transitions of C, N, and O and by Fe XVII lines."755 Above 40 tthe LETGS spectrum shows many L-shell lines of e.g.. Ne. Mg. and Si. together with lines of Fe VIII-XIII of which the Fe IX line at 171.075 is very prominent.," Above 40 the LETGS spectrum shows many L-shell lines of e.g., Ne, Mg, and Si, together with lines of Fe VIII-XIII of which the Fe IX line at 171.075 is very prominent."756 All methods applied in Sect., All methods applied in Sect.757 3.2 to the spectra of the RGS+MOS and the LETGS show temperatures of the corona of Procyon between 1-3 MK., 3.2 to the spectra of the RGS+MOS and the LETGS show temperatures of the corona of Procyon between 1–3 MK.758 No indication for à considerably higher temperature component (T> 4+ MK) is found., No indication for a considerably higher temperature component $T \ga$ 4 MK) is found.759 The total EM obtained using RGS and LETGS is about L.1(.5)«10° °., The total $EM$ obtained using RGS and LETGS is about $4.1(.5) \times 10^{50}$ $^{-3}$.760 The EAL distributioαυ] shows a smooth continuous structure without separated peak structures., The $EM$ distribution shows a smooth continuous structure without separated peak structures.761 Our results improve on those of Schmitt et al. (, Our results improve on those of Schmitt et al. (7621996) who obtain an A distribution with à maximum temperature around 1.6 MK and a cutoff beyond 6.3 MK.,1996) who obtain an $EM$ distribution with a maximum temperature around 1.6 MK and a cutoff beyond 6.3 MK.763 No significant variability of the coronal conditions took place between the observations by RGS and LETGS., No significant variability of the coronal conditions took place between the observations by RGS and LETGS.764 The abundances of C and Ν. relative to O are somewhat higher ( factor 1.5) than the values obtained in the solar photosphere (Anders Grevesse 1989).," The abundances of C and N, relative to O are somewhat higher $\sim$ factor 1.5) than the values obtained in the solar photosphere (Anders Grevesse 1989)."765 The Fe abundance is about I-1.5 « solar., The Fe abundance is about 1–1.5 $\times$ solar.766 No significance for a FIP effect. as observed in the solar corona (Feldman et al.," No significance for a FIP effect, as observed in the solar corona (Feldman et al."767 1992). is found.," 1992), is found."768 The same was concluded by Drake et al (1995). based on EUVE observations.," The same was concluded by Drake et al (1995), based on EUVE observations."769 This result is an exception of the trends found by Audard et al. (, This result is an exception of the trends found by Audard et al. (77020019) for RS CVn systems and by Giiddel et al. (,2001c) for RS CVn systems and by Güddel et al. (77120019) for solar analogs.,2001c) for solar analogs.772 These authors have found indications for the evolution from an inverse FIP effect for highly active stars - via the absence of a FIP effect in intermediately active stars - towards a normal FIP effect for less active stars., These authors have found indications for the evolution from an inverse FIP effect for highly active stars - via the absence of a FIP effect in intermediately active stars - towards a normal FIP effect for less active stars.773 Clearly. the weakly active star Procyon does not fit into this picture.," Clearly, the weakly active star Procyon does not fit into this picture."774whilst a neighbouring ray passing through a more rarefied region is not split.,whilst a neighbouring ray passing through a more rarefied region is not split.775 Once the ionization front is located along a ray. the propagation of that ray is terminated. in order to reduce the computational overhead: this is. particularly. important during the early stages of evolution. when the Strómgren sphere only involves a small fraction of the total number of SPH particles.," Once the ionization front is located along a ray, the propagation of that ray is terminated, in order to reduce the computational overhead; this is particularly important during the early stages of evolution, when the mgren sphere only involves a small fraction of the total number of SPH particles."776 Although heavy elements play an important role in determining the temperature of interstellar gas. we do not consider their chemistry here.," Although heavy elements play an important role in determining the temperature of interstellar gas, we do not consider their chemistry here."777 Instead. we assume that the composition of the gas is X=0.7 hydrogen and Y=0.3 helium. by mass: and that the helium is everywhere neutral.," Instead, we assume that the composition of the gas is $X=0.7$ hydrogen and $Y=0.3$ helium, by mass; and that the helium is everywhere neutral."778" In the neutral gas we assume that the temperature is 7,=LOK and the hydrogen is all molecular: hence the mean molecular welght is jjj,=2.35 and the isothermal sound speed is cy= 0.2kms~!."," In the neutral gas we assume that the temperature is $T_{\rm n}=10\,{\rm K}$ and the hydrogen is all molecular; hence the mean molecular weight is $\mu_{\rm n}=2.35$ and the isothermal sound speed is $c_{\rm n}=0.2\,{\rm km}\,{\rm s}^{-1}$ ."779 In the ionized gas we assume that the temperature is T;=10K and the hydrogen is all ionized: hence py=0.678 and c;=11kms7!.," In the ionized gas we assume that the temperature is $T_{\rm i}=10^4\,{\rm K}$ and the hydrogen is all ionized; hence $\mu_{\rm i}=0.678$ and $c_{\rm i}=11\,{\rm km}\,{\rm s}^{-1}$."780 In the transition region between the molecular and ionized regions. we impose a linear temperature gradient between these two limiting values.," In the transition region between the molecular and ionized regions, we impose a linear temperature gradient between these two limiting values."781 We explore three cases. (, We explore three cases. (7821) In the first case. the star lies at the centre of a spherical cloud. and the is spherically symmetric throughout its expansion. (,"i) In the first case, the star lies at the centre of a spherical cloud, and the is spherically symmetric throughout its expansion. ("78311) In the second case. the star is placed off-centre inside a spherical cloud: here the breaks out of the cloud on one side. and the remainder of the cloud ts accelerated by the rocket effect. (,"ii) In the second case, the star is placed off-centre inside a spherical cloud; here the breaks out of the cloud on one side, and the remainder of the cloud is accelerated by the rocket effect. ("784111) In the third case. the star is located outside the cloud from the outset. and à shock is driven into the cloud ahead of the ionization front (i.e. radiation driven compression).,"iii) In the third case, the star is located outside the cloud from the outset, and a shock is driven into the cloud ahead of the ionization front (i.e. radiation driven compression)."785 These cases are presented only as illustrative examples of what the code can simulate., These cases are presented only as illustrative examples of what the code can simulate.786 Detailed investigations of these phenomena will be presented in subsequent papers., Detailed investigations of these phenomena will be presented in subsequent papers.787 The paper is organized as follows., The paper is organized as follows.788 In Section 2. we discuss briefly the expansion of anregion. once the initial Strómgren sphere has formed.," In Section \ref{sec.physics} we discuss briefly the expansion of an, once the initial mgren sphere has formed."789 In Section 3 we describe in detail how we treat the propagation of ionizing radiation., In Section \ref{sec.numerical} we describe in detail how we treat the propagation of ionizing radiation.790 In Section 4 we test the algorithm on the three cases described above. and in Section 5 we discuss the results and conclude.," In Section \ref{sec.applications} we test the algorithm on the three cases described above, and in Section \ref{sec.discussion} we discuss the results and conclude."791 Consider an arbitrary density field. p(r). and suppose that there is an Ionizing star at the centre of co-ordinates.," Consider an arbitrary density field, $\rho({\bf r})$, and suppose that there is an ionizing star at the centre of co-ordinates."792" Assuming ionization equilibrium. and neglecting the diffuse radiation field. the position of the ionization front CIF). in the direction of the unit vector €. is given by R,.=R,6. where R, is obtained from Here. ni=my/X2.4x10—7& is the mass associated with one hydrogen nucleus when account is taken of the contribution from helium. zi, is the proton mass. N,,.. is the rate at which the exciting star emits Lyman continuum photons. and a, is the recombination coetficient into excited states only."," Assuming ionization equilibrium, and neglecting the diffuse radiation field, the position of the ionization front (IF), in the direction of the unit vector $\bf{\hat e}\,$, is given by ${\bf R}_{_{\rm IF}}=R_{_{\rm IF}}{\bf{\hat e}}\,$, where $R_{_{\rm IF}}$ is obtained from Here, $m=m_{\rm p}/X=2.4\times 10^{-24}\,{\rm g}$ is the mass associated with one hydrogen nucleus when account is taken of the contribution from helium, $m_{\rm p}$ is the proton mass, $\dot{\cal N}_{_{\rm LyC}}$, is the rate at which the exciting star emits Lyman continuum photons, and $\alpha_{_{\rm B}}$ is the recombination coefficient into excited states only."793 Eqn. (1)), Eqn. \ref{integral}) )794 assumes that the material inside the is fully ionized., assumes that the material inside the is fully ionized.795 Thus Eqn. (1)), Thus Eqn. \ref{integral}) )796 determines the radius Αι. at which all the ionizing photons emitted in the direction @ have been used up balancing recombinations into excited states., determines the radius $R_{_{\rm IF}}$ at which all the ionizing photons emitted in the direction $\bf{\hat e}$ have been used up balancing recombinations into excited states.797 We ignore recombinations straight into the ground state by invoking theon-the-spot approximation (Osterbrock 1974)., We ignore recombinations straight into the ground state by invoking the approximation (Osterbrock 1974).798 Strómgren (1939) was the first to show that the transition from a state of almost completely tonized material to a state of almost completely neutral material occurs 1n a very short distance compared with the dimensions of theregion., mgren (1939) was the first to show that the transition from a state of almost completely ionized material to a state of almost completely neutral material occurs in a very short distance compared with the dimensions of the.799. For example. for a spherically-symmetric expanding into a cloud of uniform density py. the radius of the ionization front is given by the Strómgren radius and the distance over which the degree of ionization changes from 90% to 106c is given by Here &=7x107env. is the mean photoionization cross section presented by a hydrogen atom to Lyman continuum photons from an OB star.," For example, for a spherically-symmetric expanding into a cloud of uniform density $\rho_{\rm n}$, the radius of the ionization front is given by the mgren radius and the distance over which the degree of ionization changes from $90\,\%$ to $10\,\%$ is given by Here $\bar\sigma=7\times10^{-18}\,{\rm cm}^2$ is the mean photoionization cross section presented by a hydrogen atom to Lyman continuum photons from an OB star."800" Because the squared sound speed in the 1onized gas inside the is et=22knyY$7. whereas the squared sound speed in the neutral material outside the is cE0.0352km? s, there is a large pressure difference betwee! the two regimes (more than three orders of magnitude). and this results in rapid expansion of the region."," Because the squared sound speed in the ionized gas inside the is $c_{\rm i}^2\simeq122\,{\rm km}^2\,{\rm s}^{-2}$, whereas the squared sound speed in the neutral material outside the is $c_{\rm n}^2\simeq 0.0352\,{\rm km}^2\,{\rm s}^{-2}$ , there is a large pressure difference between the two regimes (more than three orders of magnitude), and this results in rapid expansion of the ."801 The outward propagation of the ionization front is subsonic relative to the ionized gas (where the sound speed is c;~11 kms). but supersonic relative to the neutral gas (where the sound speed is cy02 kms7!).," The outward propagation of the ionization front is subsonic relative to the ionized gas (where the sound speed is $c_{\rm i}\sim 11\,{\rm km\,s^{-1}}$ ), but supersonic relative to the neutral gas (where the sound speed is $c_{\rm n}\sim 0.2\,{\rm km}\,{\rm s}^{-1}$ )."802 Consequently a strong shock front appears ahead the ionization front (for a detailed discussion see Kahr 1954)., Consequently a strong shock front appears ahead the ionization front (for a detailed discussion see Kahn 1954).803 Spitzer (1978) has obtained an approximate analytic solution for this phase of the evolution., Spitzer (1978) has obtained an approximate analytic solution for this phase of the evolution.804 In this solution. the radius of the ionization front is given by the density of the tonized gas by the mass of ionizedgas by and the speed at which the ionization front propagates by," In this solution, the radius of the ionization front is given by the density of the ionized gas by the mass of ionizedgas by and the speed at which the ionization front propagates by"805"cannot be ruled out as the cause of the velocity separation. (he energetics are uncomlLortably light,","cannot be ruled out as the cause of the velocity separation, the energetics are uncomfortably tight."806" Moreover. the spatio-velocity structure of the (wo components. as shown in figure 4.. shows the greatest departiures [rom the centroid velocity at around 182""00*.— the futhest projected distance Irom the central star."," Moreover, the spatio-velocity structure of the two components, as shown in figure \ref{lv}, shows the greatest departures from the centroid velocity at around $^h$ $^m$ $^s$ – the furthest projected distance from the central star."807 This does not fit the classical pattern of a shell-like expansion around a central source. and suggests that at least some component of the cloud velocily separation is unrelated (o (heir interaction with the central star We here suggest that much of the observed motion of the clouds does not arise as a direct result of the stellar cluster but is instead svstenic ie. present [rom before the clusters birth.," This does not fit the classical pattern of a shell-like expansion around a central source, and suggests that at least some component of the cloud velocity separation is unrelated to their interaction with the central star We here suggest that much of the observed motion of the clouds does not arise as a direct result of the stellar cluster but is instead systemic — i.e. present from before the cluster's birth."808 Yet the 2 kins + cloud and cloud C are clearly associated with M20 and in close spatial proximity (o one another., Yet the 2 km $^{-1}$ cloud and cloud C are clearly associated with M20 and in close spatial proximity to one another.809 We suggest (hal a scenario in which a cloud-cloud. collision triggered the formation of the central stars is highlv consistent with these observational characteristics., We suggest that a scenario in which a cloud-cloud collision triggered the formation of the central stars is highly consistent with these observational characteristics.810 In this scenario the 2 kms ! cloud and cloud C collided each other ο]. Myr ago at a relative velocity of ~7.5 km 1 or more., In this scenario the 2 km $^{-1}$ cloud and cloud C collided each other $\sim$ 1 Myr ago at a relative velocity of $\sim$ 7.5 km $^{-1}$ or more.811 The rapid collision between the (wo clouds strongly compressed the molecular gas and triggered the formation of the central star and surrounding first generation stars., The rapid collision between the two clouds strongly compressed the molecular gas and triggered the formation of the central star and surrounding first generation stars.812 In this model. the 2 km ! cloud. and cloud € must be moving in the opposite directions: the former is moving toward us and the latter is moving away [rom us.," In this model, the 2 km $^{-1}$ cloud and cloud C must be moving in the opposite directions; the former is moving toward us and the latter is moving away from us."813 As shown in Section 3.1. the 2 kins ! cloud is apparently located at the front side of M20. whereas cloud C is not in the near side. suggesting that it is located either on the far side or within M20.," As shown in Section 3.1, the 2 km $^{-1}$ cloud is apparently located at the front side of M20, whereas cloud C is not in the near side, suggesting that it is located either on the far side or within M20."814 This relative configuration is in [act consistent with the cloud-cloucl collision scenario. because we must be witnessing a moment alter (he collision occurred. Myr ago.," This relative configuration is in fact consistent with the cloud-cloud collision scenario, because we must be witnessing a moment after the collision occurred Myr ago."815 We would expect a reversed cloud location prior to the collision., We would expect a reversed cloud location prior to the collision.816 This scenario is similar to (hat cdiscussecl by Furukawaetal.(2009) and later expanded on by Ohamaetal.(2010)., This scenario is similar to that discussed by \citet{fur2009} and later expanded on by \citet{oha2010}.817. These authors Found two GAICs closely associated with the ROW 49 WIT region and its exciting cluster. the super star cluster Westerlund 2.," These authors found two GMCs closely associated with the RCW 49 HII region and its exciting cluster, the super star cluster Westerlund 2."818 The two, The two819Globular Clusters (6€) are thought to be the oldest bound stellar svstems in our Galaxy.,Globular Clusters (GC) are thought to be the oldest bound stellar systems in our Galaxy.820 Their study provides therefore valuable information. about the early Galactic. evolution., Their study provides therefore valuable information about the early Galactic evolution.821 In this respect. a major problem is that we do not. know whether what we presently observe is still representative of the initial conditions and. thus. a fossil imprint. of the formation process. or whether the initial conditions have been wiped out by a GC long evolution within the tidal fields of the Alilky Way.," In this respect, a major problem is that we do not know whether what we presently observe is still representative of the initial conditions and, thus, a fossil imprint of the formation process, or whether the initial conditions have been wiped out by a Gyr long evolution within the tidal fields of the Milky Way."822 Modelling the dynamical evolution of the Galactic Globular Cluster Svstem (GCS) is thus of ereat interest as it helps us to go back in time to the earliest stages of the cluster svstem and to disentangle the formation and evolutionary. fingerprints (see. e.g.. Cinedin Ostriker 1997. Baumgardt 1998. Vesperini 1998. Fall Zhang 2001).," Modelling the dynamical evolution of the Galactic Globular Cluster System (GCS) is thus of great interest as it helps us to go back in time to the earliest stages of the cluster system and to disentangle the formation and evolutionary fingerprints (see, e.g., Gnedin Ostriker 1997, Baumgardt 1998, Vesperini 1998, Fall Zhang 2001)."823 The GCs most vulnerable to evaporation and disruption are the low-mass clusters located at small galactocentric distance., The GCs most vulnerable to evaporation and disruption are the low-mass clusters located at small galactocentric distance.824 As a result. the evolution with time of a GCS is markedly determined by the initial distribution of the GCs in space around the Galactic centre as well as by their initial Mass spectrunmi.," As a result, the evolution with time of a GCS is markedly determined by the initial distribution of the GCs in space around the Galactic centre as well as by their initial mass spectrum."825 As for the presently. observed: spatial distribution. of the Galactic halo GC's. it is centrally concentrated with the density varving as D77 (D is the Galactocentric distance) over most of the halo (Zinn 1985).," As for the presently observed spatial distribution of the Galactic halo GCs, it is centrally concentrated with the density varying as $D^{-3.5}$ $D$ is the Galactocentric distance) over most of the halo (Zinn 1985)."826 In the inner kkpe. the distribution Uattens to something closer to an 4) τιependence.," In the inner kpc, the distribution flattens to something closer to an $D^{-2}$ dependence."827 As a result. the overall distribution is convenicntly described by a power-law with a core (see Section 2).," As a result, the overall distribution is conveniently described by a power-law with a core (see Section 2)."828 Phe observed. central [lattening probably arises from a combination of several ellects: our failure to discover some GCs in the heavily absorbed. central regions of the Galaxy. distance errors. and the real Uattening of the distribution.," The observed central flattening probably arises from a combination of several effects: our failure to discover some GCs in the heavily absorbed central regions of the Galaxy, distance errors, and the real flattening of the distribution."829 It is still unclear whether such a Hlattening is of primordial origin and rellects the initial spatial distribution of the system. or whether it has been completely determined by evolutionary. processes.," It is still unclear whether such a flattening is of primordial origin and reflects the initial spatial distribution of the system, or whether it has been completely determined by evolutionary processes."830 “The latter are especially effective at small galactocentric distances where the GC relaxation time is short. causing the cisruοι of some GC's and the partial evaporation of some others.," The latter are especially effective at small galactocentric distances where the GC relaxation time is short, causing the disruption of some GCs and the partial evaporation of some others."831 More generally. as far as," More generally, as far as"832" where Ais the area of a cell face (see figure 1)) aud the £s are fluxes defined as To obtain the deusitics at the iterfaces (97,4) straight averages are computed of centered densities adjacent to the interface.",as where $A$ is the area of a cell face (see figure \ref{fig:cell}) ) and the $F$ s are fluxes defined as To obtain the densities at the interfaces $\rho_{i\pm 1/2}^n$ ) straight averages are computed of centered densities adjacent to the interface.833 The two centered subscripts have intentionally con. onütted In the expressions for the fluxes aud areas to reduce equation leneth (e.g. { was left off mt 71/2 was ΚΟ))., The two centered subscripts have intentionally been omitted in the expressions for the fluxes and areas to reduce equation length (e.g. $i$ was left off but $i+1/2$ was kept).834 Note that the densities iji the fluxes aro af tine 5 and not »|1/2. so «mr solution algorithui is sraiehttorwardly explicit. as ds true for many 1D calculations.," Note that the densities in the fluxes are at time $n$ and not $n+1/2$, so our solution algorithm is straightforwardly explicit, as is true for many 1D calculations."835 It is uot possible to properly time ceuter all terms without introducing a more conrplex inyplicit or imulti-«ep explicit algorithin., It is not possible to properly time center all terms without introducing a more complex implicit or multi-step explicit algorithm.836 Wih this expression for the fluxes. we can cirectly solve equation (10)) for the ceusity at the new time step.," With this expression for the fluxes, we can directly solve equation \ref{eq:mass-cons-fv}) ) for the density at the new time step."837 The final piece needed o coniplete the description is the calculation of the erid velocity., The final piece needed to complete the description is the calculation of the grid velocity.838 For a spherical shell to have costant mass. the net flow of 1iass into and out of hat spherical shell ust be zero.," For a spherical shell to have constant mass, the net flow of mass into and out of that spherical shell must be zero."839 Suwine up all he fluxes iuto and out of the individua| horizoutal cells iu a spherical shell. substituting ¢quation (11 )) iu for the outer radial flux (at ὁ| 2) and setine the result equal to zero we arrive at the equation. Solving for the outer erid velocity. (wetUrilye|l2 produces an equation for calculating the new erid velocity. The inner radial flux. FYy. is depenudeu ou the erid velocity at the inner interface.," Summing up all the fluxes into and out of the individual horizontal cells in a spherical shell, substituting equation \ref{eq:outer-flux}) ) in for the outer radial flux (at $i+1/2$ ) and setting the result equal to zero we arrive at the equation, Solving for the outer grid velocity, $v_{0r,i+1/2}^{n+1/2}$, produces an equation for calculating the new grid velocity, The inner radial flux, $F_{i-1/2}^{n+1/2}$, is dependent on the grid velocity at the inner interface."840 A the first radial zone boundary next to the ei core we impose both a zero radial velocity anc erid velocity., At the first radial zone boundary next to the rigid core we impose both a zero radial velocity and grid velocity.841 Thus. equation (15)) can be solve recursively frou. the model iterior boundary to the surface to determine the erid velocity at al interfaces.," Thus, equation \ref{eq:rad-grid-vel}) ) can be solved recursively from the model interior boundary to the surface to determine the grid velocity at all interfaces."842 The initial model for our adiabatic simulations is generated by requiring that it be iu livdrostatic equilibrimu., The initial model for our adiabatic simulations is generated by requiring that it be in hydrostatic equilibrium.843 When this constraint is applied to the conservation equations the ouly terms that remain are the pressure and eravitv terms in the radial momentum couscrvation equation., When this constraint is applied to the conservation equations the only terms that remain are the pressure and gravity terms in the radial momentum conservation equation.844 Iu particular. there are no terms left im the internal energy conservation equation. and thus no equation to solve for the energv structure.," In particular, there are no terms left in the internal energy conservation equation, and thus no equation to solve for the energy structure."845" To provide this information. an enerev profile παν. generaed from aotier stelar modeling code ROTORC (DeuprecL990) Hand eicrgies were interpolaed in loe(AL.) ο cel cenutcYs,"," To provide this information, an energy profile was generated from another stellar modeling code ROTORC \citep{Deupree-1990} and energies were interpolated in $\log(M_r)$ to cell centers."846 Once we impose he energv disributiol. we can sinultaneouslv solve the racial wdrostatic ecuilibriuui fuite differeice equation a1 the equation of state for the pressure aud ceusivo structure of the model given he spacing of he indepevent variahle AZ.," Once we impose the energy distribution, we can simultaneously solve the radial hydrostatic equilibrium finite difference equation and the equation of state for the pressure and density structure of the model given the spacing of the independent variable $M_r$."847 The Lacjus is deernüned from the vohune required to produce the calculated density from the iass of the shell., The radius is determined from the volume required to produce the calculated density from the mass of the shell.848 No convecive model is incbluded iu the starting model because RR Lyrae do not have extensive convective regions to affect the structure., No convective model is included in the starting model because RR Lyrae do not have extensive convective regions to affect the structure.849 To induce pulsation a racial velocity profile from the linear. noun-adiabatic. radial," To induce pulsation a radial velocity profile from the linear, non-adiabatic, radial"850"ray and cosmic ray measurements, is that it is practically insensitive to the complications caused by the nonlinear evolution of the cosmic density field.","ray and cosmic ray measurements, is that it is practically insensitive to the complications caused by the nonlinear evolution of the cosmic density field."851" As our analysis shows, for any realistic structure formation scenario the CMB bounds on annihilating DM arise solely around the redshifts of z~1000, while the contribution from lower redshift cosmic structures is completely negligible."," As our analysis shows, for any realistic structure formation scenario the CMB bounds on annihilating DM arise solely around the redshifts of $z\sim 1000$, while the contribution from lower redshift cosmic structures is completely negligible."852 CMB constraints on annihilating DM have been obtained in several earlier studies: e.g. ??7?????..," CMB constraints on annihilating DM have been obtained in several earlier studies: e.g. \citet{2005PhRvD..72b3508P,2006MNRAS.369.1719M,2006PhRvD..74j3519Z,2009PhRvD..80b3505G,2009PhRvD..80d3526S,2009JCAP...10..009C,2009A&A...505..999H,2010PThPh.123..853K}."853" Even though most of these analyses have assumed a simple ‘on the spot’ approximation for the energy more recent studies (??) followed the energydeposition}, transport problem including various energy-loss mechanisms in a more realistic way."," Even though most of these analyses have assumed a simple `on the spot' approximation for the energy, more recent studies \citep{2009PhRvD..80d3526S,2009A&A...505..999H} followed the energy transport problem including various energy-loss mechanisms in a more realistic way."854" Compared to the analysis of ?,, which partially relies on the previously derived ‘on the spot’ results of ?,, in this paper we perform a more elaborate treatment for the energy deposition joined to the analysis of the CMB data via Markov chain Monte Carlo calculations that incorporate the most recent WMAP likelihood code."," Compared to the analysis of \citet{2009PhRvD..80d3526S}, which partially relies on the previously derived `on the spot' results of \citet{2009PhRvD..80b3505G}, in this paper we perform a more elaborate treatment for the energy deposition joined to the analysis of the CMB data via Markov chain Monte Carlo calculations that incorporate the most recent WMAP likelihood code."855" Also, we make an attempt to unify the results from various annihilation channels and provide a simple and rather generic fitting formula for calculating CMB constraints on annihilation cross section (cv) for a broad range of annihilating DM models."," Also, we make an attempt to unify the results from various annihilation channels and provide a simple and rather generic fitting formula for calculating CMB constraints on annihilation cross section $\cs$ for a broad range of annihilating DM models."856 Our paper is organized as follows., Our paper is organized as follows.857 In Section 2 we give a brief description of the energy input from DM annihilation and provide a simple treatment for its propagation., In Section 2 we give a brief description of the energy input from DM annihilation and provide a simple treatment for its propagation.858 The effect on CMB temperature and polarization fluctuations is investigated in Section 3., The effect on CMB temperature and polarization fluctuations is investigated in Section 3.859 Section 4 presents our main results about current and future CMB constraints., Section 4 presents our main results about current and future CMB constraints.860 Our summary is given in Section 5., Our summary is given in Section 5.861" As in ??,, we treat our input signals from DM annihilation in an as model independent a way as possible."," As in \citet{2009NuPhB.813....1C,2011JCAP...03..051C}, we treat our input signals from DM annihilation in an as model independent a way as possible."862" In? the two-particle annihilation channels to all Standard Model (SM) particles were considered: leptons, quarks, photons, gluons, weak-interaction gauge bosons, Higgs boson, and neutrinos."," In \citet{2011JCAP...03..051C} the two-particle annihilation channels to all Standard Model (SM) particles were considered: leptons, quarks, photons, gluons, weak-interaction gauge bosons, Higgs boson, and neutrinos."863" In addition, annihilations to four leptons via an intermediate new boson V were considered."," In addition, annihilations to four leptons via an intermediate new boson $V$ were considered."864 Such a treatment can be considered as model independent since realistic models can always be decomposed into these basic channels where the particular branching ratios between the channels are given by the underlying theoretical particle physics model., Such a treatment can be considered as model independent since realistic models can always be decomposed into these basic channels where the particular branching ratios between the channels are given by the underlying theoretical particle physics model.865" Since in our work we focus on DM particle masses below 100 GeV, out of allof the above channels the following remain: DM DM — SM SM, where SM={e, ,7,9,¢,0,7,9}3; plus 4-lepton channels via V."," Since in our work we focus on DM particle masses below $100$ GeV, out of allof the above channels the following remain: DM DM $\rightarrow$ SM $\overline{{\rm SM}}$, where $=\{e,\mu,\tau,q,c,b,\gamma,g$ ; plus 4-lepton channels via $V$."866" Here q denotes the light quarks u, d, and s."," Here $q$ denotes the light quarks $u$, $d$, and $s$ ."867" Because the masses of interest in this work are mostly below the masses of the electroweak gauge bosons, W- and Z, those channels are left out."," Because the masses of interest in this work are mostly below the masses of the electroweak gauge bosons, $W^{\pm}$ and $Z$, those channels are left out."868 For the same reason we also do not need to distinguish between left- and right-handed particles., For the same reason we also do not need to distinguish between left- and right-handed particles.869" Also, the neutrino channels in this case provide only trivial output; i.e. 10096 of the energy is carried away by neutrinos, which escape freely at the redshifts of interest, so are not treated any further."," Also, the neutrino channels in this case provide only trivial output; i.e. $100\%$ of the energy is carried away by neutrinos, which escape freely at the redshifts of interest, so are not treated any further."870" Even though the channels y and g are included in our model-independent approach, these are strongly suppressed for realistic models since DM should not carry color or interact electromagnetically."," Even though the channels $\gamma$ and $g$ are included in our model-independent approach, these are strongly suppressed for realistic models since DM should not carry color or interact electromagnetically."871" For all channels, the spectra of the emerging stable particles, e, p, y, v, after treatment of several decays, parton showers, and hadronization were calculated using PYTHIA Monte (?).."," For all channels, the spectra of the emerging stable particles, $e$, $p$, $\gamma$, $\nu$, after treatment of several decays, parton showers, and hadronization were calculated using PYTHIA Monte \citep{2008CoPhC.178..852S}."872 All input spectra are downloadable from[PPPCADMID., All input spectra are downloadable from.873"html.. For more details, and in particular for a discussion on the level of possible uncertainties, we refer the reader to ?.."," For more details, and in particular for a discussion on the level of possible uncertainties, we refer the reader to \citet{2011JCAP...03..051C}."874" Among the stable output particles, neutrinos propagate freely at the redshifts of interest, while e* immediately interact with the ubiquitous CMB photons and upscatter those to the gamma-ray energy range via the inverse Compton (IC) mechanism."," Among the stable output particles, neutrinos propagate freely at the redshifts of interest, while $e^{\pm}$ immediately interact with the ubiquitous CMB photons and upscatter those to the gamma-ray energy range via the inverse Compton (IC) mechanism."875 The total output energy in hadrons (p and d) is typically quite negligible., The total output energy in hadrons $p$ and $d$ ) is typically quite negligible.876" Only in gluon and quark channels does it reach up to ~ 1596, and that only for thehighestDM particle masses considered in this paper."," Only in gluon and quark channels does it reach up to $\sim 15\%$ , and that only for thehighestDM particle masses considered in this paper."877 We therefore did not model this component in detail., We therefore did not model this component in detail.878" However, in Sect."," However, in Sect."879 B] we give some, \ref{sec:protons} we give some880"kinematics and stellar populations (?) obtained using exactly the same data analysis technique, i.e. the full spectral fitting FLAMES-LR04 spectra having similar although slightly wider wavelength range but twice lower spectral resolution.","kinematics and stellar populations \citep{CCB08} obtained using exactly the same data analysis technique, i.e. the full spectral fitting FLAMES-LR04 spectra having similar although slightly wider wavelength range but twice lower spectral resolution."881 The velocity dispersion measurements agree remarkably well for all three objects., The velocity dispersion measurements agree remarkably well for all three objects.882" The metallicity measurements agree within a few hundredths dex for UCD 3 and UCD 4, however being discrepant by ~0.4 dex for UCD 2."," The metallicity measurements agree within a few hundredths dex for UCD 3 and UCD 4, however being discrepant by $\sim8830.4$ dex for UCD 2."884" For this object, the age estimations also differ: intermediate in ? and old in our present study."," For this object, the age estimations also differ: intermediate in \citet{CCB08} and old in our present study."885" We notice, however, that the discrepancy of the UCD 2 age and metallicity measurements in ? and our present study follow the age-metallicity degeneracy."," We notice, however, that the discrepancy of the UCD 2 age and metallicity measurements in \citet{CCB08} and our present study follow the age–metallicity degeneracy."886 Given quite poor data quality and large size of stellar population confidence levels in ? we conclude that the discrepancy between the measurements can be explained by statistical effects., Given quite poor data quality and large size of stellar population confidence levels in \citet{CCB08} we conclude that the discrepancy between the measurements can be explained by statistical effects.887 Age determinations for UCD 3 and UCD 4 agree between the two studies within 2c of their statistical uncertainties., Age determinations for UCD 3 and UCD 4 agree between the two studies within $\sigma$ of their statistical uncertainties.888" It is worth mentioning, that the metallicity estimates for all three UCDs presented in ? are systematically lower by 0.20...0.25 dex compared to our present measurements."," It is worth mentioning, that the metallicity estimates for all three UCDs presented in \citet{MHIJ06} are systematically lower by $\dots$ 0.25 dex compared to our present measurements."889 In Fig 7 (black data points) we present the comparison of published velocity dispersion measurements for 19 UCDs (?) with those obtained in our study with the full-spectral fitting technique.," In Fig \ref{figsigsig} (black data points) we present the comparison of published velocity dispersion measurements for 19 UCDs \citep{Mieske+08}890 with those obtained in our study with the full-spectral fitting technique."891" Although the general trend agrees, the measurements for individual objects are often notably discrepant."," Although the general trend agrees, the measurements for individual objects are often notably discrepant."892 The reasons for the discrepancy are: (1) template mismatch during cross-correlation due to the metallicity difference between UCDs and w Cen giant stars served as templates; (2) slightly different wavelength ranges used for the data analysis (inclusion of the Mgb triplet region in our study); (3) our correction for the contamination of UCD spectra by the NGC 1399 halo which ? did not apply., The reasons for the discrepancy are: (1) template mismatch during cross-correlation due to the metallicity difference between UCDs and $\omega$ Cen giant stars served as templates; (2) slightly different wavelength ranges used for the data analysis (inclusion of the $b$ triplet region in our study); (3) our correction for the contamination of UCD spectra by the NGC 1399 halo which \citet{Mieske+08} did not apply.893" The metallicity difference between the spectra and templates used to analyse them may lead to biased velocity dispersion measurements (see Section 1.3.1 in ?)) at least if the analysis is done in the pixel space, thus affecting both full spectral fitting andFXCOR cross-correlation measurements."," The metallicity difference between the spectra and templates used to analyse them may lead to biased velocity dispersion measurements (see Section 1.3.1 in \citealp{Chilingarian06}) ) at least if the analysis is done in the pixel space, thus affecting both full spectral fitting and cross-correlation measurements."894" The low metallicity of a template star resulting in shallower absorption lines may be compensated by decreasing the velocity dispersion, i.e. smearing absorption lines to a lower degree than it should be in order to match the line depth in the target spectrum being analysed."," The low metallicity of a template star resulting in shallower absorption lines may be compensated by decreasing the velocity dispersion, i.e. smearing absorption lines to a lower degree than it should be in order to match the line depth in the target spectrum being analysed."895" Then we would expect to see the correlation between UCD metallicities and differences of velocity dispersion measurements in ? and our present study, which we do not detect at a statistically significant level."," Then we would expect to see the correlation between UCD metallicities and differences of velocity dispersion measurements in \citet{Mieske+08} and our present study, which we do not detect at a statistically significant level."896 This can be explained because the described degeneracy between metallicity and velocity dispersion mostly affects the data for targets with velocity dispersions similar to or lower than the instrumental spectral resolution., This can be explained because the described degeneracy between metallicity and velocity dispersion mostly affects the data for targets with velocity dispersions similar to or lower than the instrumental spectral resolution.897" In our case, most c targets have low metallicities well corresponding to those of the w Cen template stars."," In our case, most $\sigma$ targets have low metallicities well corresponding to those of the $\omega$ Cen template stars."898" On the other hand, massive metal-rich UCDs with relatively high velocity dispersions are bright,hence their spectra have good signal-to-noise ratios reducing the degeneracy effects."," On the other hand, massive metal-rich UCDs with relatively high velocity dispersions are bright,hence their spectra have good signal-to-noise ratios reducing the degeneracy effects."899 ? had to exclude the Mgb triplet region from their analysis as it seemed to bias the velocity dispersion measurements obtained using the cross-correlation technique., \citet{Mieske+08} had to exclude the $b$ triplet region from their analysis as it seemed to bias the velocity dispersion measurements obtained using the cross-correlation technique.900" However, Mgb is the most prominent spectral feature in the wavelength range of our spectra, thus containing a large fraction of spectral information."," However, $b$ is the most prominent spectral feature in the wavelength range of our spectra, thus containing a large fraction of spectral information."901 We performed the spectral fitting in the wavelength range Acestframe>5200 aand compared the measurements of velocity dispersion with those obtained from the fitting in the entire available wavelength range aimed at checking whether the full spectral fitting technique also suffers from similar biases.," We performed the spectral fitting in the wavelength range $\lambda_{\mathrm{restframe}} >9025200$ and compared the measurements of velocity dispersion with those obtained from the fitting in the entire available wavelength range aimed at checking whether the full spectral fitting technique also suffers from similar biases."903" The values turned to be consistent within their uncertainties, however, being almost half as precise in case of the truncated wavelength range."," The values turned to be consistent within their uncertainties, however, being almost half as precise in case of the truncated wavelength range."904" Therefore, we conclude that the biases of velocity dispersion measurements obtained by cross-correlation including the Mgb triplet in the wavelength range probably originate from the template mismatch when using stellar spectra as references, which is minimized in our case by selecting the best-matching SSP from the grid of stellar population models."," Therefore, we conclude that the biases of velocity dispersion measurements obtained by cross-correlation including the $b$ triplet in the wavelength range probably originate from the template mismatch when using stellar spectra as references, which is minimized in our case by selecting the best-matching SSP from the grid of stellar population models."905" As far as these biases affect only measurements of low velocity dispersions in objects like globular clusters made on relatively high resolution spectra and do not seem to show up in studies of relatively massive galaxies, we suppose that this mismatch between the spectra of individual stars and unresolved stellar populations originates from subtle absorption-line features and becomes important only at high spectral resolution."," As far as these biases affect only measurements of low velocity dispersions in objects like globular clusters made on relatively high resolution spectra and do not seem to show up in studies of relatively massive galaxies, we suppose that this mismatch between the spectra of individual stars and unresolved stellar populations originates from subtle absorption-line features and becomes important only at high spectral resolution."906" Finally, we repeated the velocity dispersion measurements of ? using the task inIRAF, but now applied to the UCD spectra corrected for the contamination of NGC 1399."," Finally, we repeated the velocity dispersion measurements of \citet{Mieske+08} using the task in, but now applied to the UCD spectra corrected for the contamination of NGC 1399."907" Using the same stellar templates of w Cen as in ?,, the discrepancy to the results obtained with the technique has dramatically decreased which is clearly seen in Fig 7 (red data points)."," Using the same stellar templates of $\omega$ Cen as in \citet{Mieske+08}, the discrepancy to the results obtained with the technique has dramatically decreased which is clearly seen in Fig \ref{figsigsig} (red data points)."908" We obtain à good agreement between the two datasets, with only a small residual systematic offset in the sense that dispersions are slightly larger thanNBURST dispersions for low dispersion values."," We obtain a good agreement between the two datasets, with only a small residual systematic offset in the sense that dispersions are slightly larger than dispersions for low dispersion values."909" Thus, we conclude that the contamination of the spectra by the host galaxy halo is the main reason for biases of the estimated kinematical"," Thus, we conclude that the contamination of the spectra by the host galaxy halo is the main reason for biases of the estimated kinematical"910observations were made in the compact array configuration of the SMA. where the projected shortest and longest baselines were ~ [4 m and ~ 69 m CIO and 53 kA) respectively.,"observations were made in the compact array configuration of the SMA, where the projected shortest and longest baselines were $\sim$ 14 m and $\sim$ 69 m (10 and 53 $\lambda$ ) respectively."911 We used Uranus for both bandpass calibration and flux calibration., We used Uranus for both bandpass calibration and flux calibration.912 Amplitude and phase calibrations were done with the quasar 1751+096 and the quasar 1743-038 was observed to verify the quality of phase referencing from 17514096., Amplitude and phase calibrations were done with the quasar 1751+096 and the quasar 1743-038 was observed to verify the quality of phase referencing from 1751+096.913 The visibility data were calibrated using the MIR package and the maps were generated and CLEANed using the NRAO AIPS package., The visibility data were calibrated using the MIR package and the maps were generated and CLEANed using the NRAO AIPS package.914 Even after calibration the continuum visibility data of the reference quasar 1743-038 showed evidence for phase decorrelation at longer baselines., Even after calibration the continuum visibility data of the reference quasar 1743-038 showed evidence for phase decorrelation at longer baselines.915 The degree of decorrelation was similar for 1743-038 and MWC 297., The degree of decorrelation was similar for 1743-038 and MWC 297.916 We self calibrated the visibility data for both 1743-038 and MWC 297 at a time interval of 4 min. which significantly reduced the phase decorrelation.," We self calibrated the visibility data for both 1743-038 and MWC 297 at a time interval of 4 min, which significantly reduced the phase decorrelation."917 The vector averaged visibility amplitude plotted against the av distance for MWC 297 and the quasar 1743-038 after self calibration is shown in Fig. 1.., The vector averaged visibility amplitude plotted against the $uv$ distance for MWC 297 and the quasar 1743-038 after self calibration is shown in Fig. \ref{vis}.918 We imaged the self calibrated continuum visibilities of MWC 297., We imaged the self calibrated continuum visibilities of MWC 297.919 The resultant size ofthe synthesized beam was 3.71] « 2.97 with uniform weighting (PA — 45. )., The resultant size ofthe synthesized beam was $^{\prime \prime}$ 11 $ \times$ $^{\prime \prime}$ 97 with uniform weighting (PA $\sim$ 45 $^{\degr}$ ).920 The contour map of the observed continuum emission at 1.3 mm is shown in Fig. 2.., The contour map of the observed continuum emission at 1.3 mm is shown in Fig. \ref{map}. .921 We expect the maximum uncertainty in source positions to be <0.ff3 based on the positions of quasars mapped in our SMA observations., We expect the maximum uncertainty in source positions to be $\leqslant$ $0.^{\arcsec}3$ based on the positions of quasars mapped in our SMA observations.922 Uncertainties in the absolute flux is estimated to be20%., Uncertainties in the absolute flux is estimated to be.923. Compact continuum emission with a total measured flux density of 300 mJy is detected towards MWC 297 at 1.3 mm., Compact continuum emission with a total measured flux density of 300 mJy is detected towards MWC 297 at 1.3 mm.924 The angular separation between the 1.3 mm continuum peak and the stellar position is less than 0. 3., The angular separation between the 1.3 mm continuum peak and the stellar position is less than $^{\arcsec}$ 3.925 Therefore. we assume that the compact continuum source is centered upon the star.," Therefore, we assume that the compact continuum source is centered upon the star."926" A two-dimensional Gaussian fit to the continuum map of MWC 297 yields an observed size of 3.” 14 « 3.702 and a deconvolved size of 0.76470, « OLINOF at PA = + ."," A two-dimensional Gaussian fit to the continuum map of MWC 297 yields an observed size of $^{\prime927 \prime}$ 14 $ \times$ $^{\prime \prime}$ 02 and a deconvolved size of $^{\prime \prime}$ $^{+0.^{\prime \prime}10}_{-0.^{\prime928 \prime}64}$ $\times$ $^{\prime \prime}$ $^{+0.^{\prime929 \prime}60}_{-0.^{\prime \prime}11}$ at PA = $^{\degr}$ $\pm$ $^{\degr}$."930 Allowing for the possibility of the phase decorrelation in our observations affecting the observed size. we take the longer dimension of the deconvolved size as an upper limit to the source size.," Allowing for the possibility of the phase decorrelation in our observations affecting the observed size, we take the longer dimension of the deconvolved size as an upper limit to the source size."931 At the distance of MWC 297. it gives a source radius of 80 AU.," At the distance of MWC 297, it gives a source radius of 80 AU."932 MWC 297 is known to have an ionized wind associated with it (Malbetetal.2007:Drew1997).," MWC 297 is known to have an ionized wind associated with it \citep{malbet07,drew97}."933". The flux densities at 3.6 em and 6 cm. measured towards MWC 297 with the VLA (Skinneretal.1993).. give a spectral index of 0.6 (E, x 1°). appropriate for free-free emission from an optically thick. tronized wind."," The flux densities at 3.6 cm and 6 cm, measured towards MWC 297 with the VLA \citep{skin93}, give a spectral index of 0.6 $_{\nu}$ $\propto$ $\nu^{0.6}$ ), appropriate for free-free emission from an optically thick, ionized wind."934 Assuming that this emission continues to the mm wavelengths with the same spectral index. we subtracted the possible contribution due to the free-free emission from the total observed flux at 1.3 mm and obtained a flux density of 200 mJy as due to dust emission.," Assuming that this emission continues to the mm wavelengths with the same spectral index, we subtracted the possible contribution due to the free-free emission from the total observed flux at 1.3 mm and obtained a flux density of 200 mJy as due to dust emission."935" If the dust emission is optically thin at mm wavelengths. assuming a gas-to-dust mass ratio of 100 and a dust opacity per unit mass of dust plus gas &,, = [5272] eng? (Cesaronietal.2007:Kramer1998)... we obtain a total mass of 0.07 M... for the circumstellar material (gas+dust) associated with MWC 297."," If the dust emission is optically thin at mm wavelengths, assuming a gas-to-dust mass ratio of 100 and a dust opacity per unit mass of dust plus gas $\kappa_{\nu}$ = $\left[ \frac{\nu (GHz)}{230.6} \right936]^{\beta}$ $cm^2g^{-1}$ \citep{cesaroni07,kramer98}, we obtain a total mass of 0.07 $_{\odot}$ for the circumstellar material (gas+dust) associated with MWC 297."937 A dust temperature of 100 K is assumed in this calculation., A dust temperature of 100 K is assumed in this calculation.938 From the mass and the upper limit to the source radius. we computed the optical extinctions along the line of sight through the continuum source to the central star to probe the geometry of the circumstellar material.," From the mass and the upper limit to the source radius, we computed the optical extinctions along the line of sight through the continuum source to the central star to probe the geometry of the circumstellar material."939 If. the circumstellar material is. distributed in à spherical envelope of uniform density. the optical. V-band extinction would be 107 mag.," If the circumstellar material is distributed in a spherical envelope of uniform density, the optical V-band extinction would be $\ge$ $^4$ mag."940 However. the observed extinction Ay is only 8 mag (Drewetal.1997).," However, the observed extinction $A_V$ is only 8 mag \citep{drew97}."941. This strongly suggests that the circumstellar dust is likely distributed in a comparatively flattened and inclined. morphology around MWC 297. perhaps in the form of a disk.," This strongly suggests that the circumstellar dust is likely distributed in a comparatively flattened and inclined morphology around MWC 297, perhaps in the form of a disk."942 In Fig., In Fig.943 3 we present the spectral energy distribution (SED) of MWC 297 at submm. mm andem wavelengths.," \ref{sed} we present the spectral energy distribution (SED) of MWC 297 at submm, mm andcm wavelengths."944 We fit the observed points with a combination of free-free emission arising in an ionized wind and optically thin dust emission from a circumstellar disk of mass 0.07 ..., We fit the observed points with a combination of free-free emission arising in an ionized wind and optically thin dust emission from a circumstellar disk of mass 0.07 $_{\odot}$.945 The best fit is obtained when the dust opacity power-law exponent ./ has a value between 0.1 and 0.3., The best fit is obtained when the dust opacity power-law exponent $\beta$ has a value between 0.1 and 0.3.946 This ts much lower than that observed for interstellar grains (./ 2 2) and the representative value used for circumstellar disks (./ = 1)., This is much lower than that observed for interstellar grains $\beta$ = 2) and the representative value used for circumstellar disks $\beta$ = 1).947 The low value of } 1s often interpreted as arising due to the average size of the emitting dust grains being relatively larger., The low value of $\beta$ is often interpreted as arising due to the average size of the emitting dust grains being relatively larger.948 Larger grain size in the circumstellar environment around MWC 297 has also been reported independently from studies on the wavelength dependence of extinction. towards the star at the optical wavelengths (Gorti&Bhatt 1993).., Larger grain size in the circumstellar environment around MWC 297 has also been reported independently from studies on the wavelength dependence of extinction towards the star at the optical wavelengths \citep{gortibhatt93}. .949" This argues for possible graingrowth in the optically thin circumstellar material around MWC 297,", This argues for possible graingrowth in the optically thin circumstellar material around MWC 297.950 In order to construct the SED and to the value of ./ we have used large beam (~ ) determinesingle dish measurements. at submm wavelengths with JCMT," In order to construct the SED and to determine the value of $\beta$ , we have used large beam $\sim$ $^{\arcsec}$ ) single dish measurements at submm wavelengths with JCMT"951 , 952source being a binary companion or to have formed from the same cloud core as MWC 297 (Li ct al.,source being a binary companion or to have formed from the same cloud core as MWC 297 (Li et al.953 1991)., 1994).954 Accurate photometry of the newly fouud object is challenging as it is found close to the bright host star., Accurate photometry of the newly found object is challenging as it is found close to the bright host star.955 Usiug the 2MAÀSS Παπά imaenitude of NWC 297. Wwe caibrated our zeropoiut by usine open-loop nuages of AIWC 297.," Using the 2MASS $H$ -band magnitude of MWC 297, we calibrated our zeropoint by using open-loop images of MWC 297."956 PSF photometry ou the companion was perforued by fitting Gaussians., PSF photometry on the companion was performed by fitting Gaussians.957 In addition. we performed apertire photometry.," In addition, we performed aperture photometry."958 The resulting £7-haxd magnitudes agree arly well with each other. aud we fiud an Z7-baud imaenitude dierence of ATF = 8.5 + 0.25 mae.," The resulting $H$ -band magnitudes agree fairly well with each other, and we find an $H$ -band magnitude difference of $\Delta H$ = 8.5 $\pm$ 0.25 mag."959 Using the 2QATASS inaguitude of AWC 297. we fine II = 12.9 4 0.25 mae for he source.," Using the 2MASS magnitude of MWC 297, we find $H$ = 12.9 $\pm$ 0.25 mag for the source."960 To investigae whether our discovered ob.ject at a PA of ccould be resx»usible for the N-rav flaring of MNC 297. we turn to the relevant archival X-ray satelite data.," To investigate whether our discovered object at a PA of could be responsible for the X-ray flaring of MWC 297, we turn to the relevant archival X-ray satellite data."961 The N-rav inage of the region near MWC 297 is shown in the Ieft panel of Fig. 2..," The X-ray image of the region near MWC 297 is shown in the left panel of Fig. \ref{f_xray},"962 where we note four objects., where we note four objects.963 The positions of these objects are listed iu Table 1.., The positions of these objects are listed in Table \ref{t_chan}. .964 The two X-ray sources with a separation of & aat PA of ~ aare an exact iatcli to our ZJ-baud inagiug of AIWC 297 (source 1) and our newly discovered source (source 2)., The two X-ray sources with a separation of $\simeq$ at PA of $\simeq$ are an exact match to our $H$ -band imaging of MWC 297 (source 1) and our newly discovered source (source 2).965 Note that if Source 1 and 2 form a binary system. the large period implied by the 850 AU separation cans that proper motion will not be significant. and source 2 will be at the same position in theA and images.," Note that if Source 1 and 2 form a binary system, the large period – implied by the 850 AU separation – means that proper motion will not be significant, and source 2 will be at the same position in the and images."966 We cannot exclude the possibility that the remaimine A-rav cussion from Source 1. which is cousistent with ADWC 2977's position. is vet due to one or more other conipauions tha reniadu uuresolved in the nuage.," We cannot exclude the possibility that the remaining X-ray emission from Source 1, which is consistent with MWC 297's position, is yet due to one or more other companions that remain unresolved in the image."967" Also visible in t10 1HAGE aro SOUPCCS 3 and 4. at roughly from, AIWC 297."," Also visible in the image are sources 3 and 4, at roughly from MWC 297."968 Source counts were extracted using a radius of 3 pixes for sources l and 2. auc a radius of 5 jxels for sources 3 aud d.," Source counts were extracted using a radius of 3 pixels for sources 1 and 2, and a radius of 5 pixels for sources 3 and 4."969 The total 0.3.10. keV. source counts are listed in Table 1.., The total 0.3–10 keV source counts are listed in Table \ref{t_chan}.970 Note tha the sum of source counts from sources 3 and Lis about του times ereater han that for sources 1 aud 2., Note that the sum of source counts from sources 3 and 4 is about three times greater than that for sources 1 and 2.971 These cata raise he question of whether the low spatial resolution of nav have led to a uisiclentification of he reported N-arax daring of MAVC 297., These data raise the question of whether the low spatial resolution of may have led to a misidentification of the reported X-ray flaring of MWC 297.972 We therefore tu to the flaring data., We therefore turn to the flaring data.973 Fig., Fig.974 2) Gight) shows the combines SIS image of the flaring data on AIWC 297 ou the same scale as he nage., \ref{f_xray}~ (right) shows the combined SIS image of the flaring data on MWC 297 – on the same scale as the image.975 We have indicated the positions of sources 1. 3. and d.," We have indicated the positions of sources 1, 3, and 4."976 The cices in this nuage ive radi of191. indicatiug fre combined aud position uncertainty.," The circles in this image have radii of, indicating the combined and position uncertainty."977 It is inunnediatelv apparent youn Fig., It is immediately apparent from Fig.978 2 (ight) that the sotrees are confused in the image., \ref{f_xray}~ (right) that the sources are confused in the image.979 Moreover. tje muaproved astrometry xovided by the corrections of Gothelf e al. (," Moreover, the improved astrometry provided by the corrections of Gotthelf et al. ("9802000) shows hat t16 peak of the observed emission iu the PSF Is inconsistent with the position of sources la xd 2.,2000) shows that the peak of the observed emission in the PSF is inconsistent with the position of sources 1 and 2.981 TIustead. he peak is consistent with the positiois of source 3Pa Or l.," Instead, the peak is consistent with the positions of source 3 or 4."982 This strongly sugeestsCoco that the origin of the firing vchaviour is due to source 3 or Ll., This strongly suggests that the origin of the flaring behaviour is due to source 3 or 4.983 Their yositions dine up with 241ASS point sources. with /7 = 11.3 and 9.1 for sources 3 and Ld respectively.," Their positions line up with 2MASS point sources, with $H$ = 11.3 and 9.4 for sources 3 and 4 respectively."984 Caven the fac that Source thas a hager (PFIv) IR excess than SotPCC 3. the flaring is most Likely due to source [.," Given the fact that Source 4 has a larger $(H-K)$ IR excess than Source 3, the X-ray flaring is most likely due to source 4."985 Tu any case. at the distance of AIWC 297. these ZI-baud magnitudes are consistent with a T Tauri nature.," In any case, at the distance of MWC 297, these $H$ -band magnitudes are consistent with a T Tauri nature."986 We have presented high resolution AO-NIR and Nora Huaging on MWC 297., We have presented high resolution AO-NIR and X-ray imaging on MWC 297.987 Caven the carly spectral type of the object (B1.5). the reported X-rav flaring of this carly Herbie Be star has been diffieult to unudoerstaud.," Given the early spectral type of the object (B1.5), the reported X-ray flaring of this early Herbig Be star has been difficult to understand."988 UsingCILANDRA. we lave resolved the N-rayv cussion from. objects suvoundiug MW 297 and found that t1C brightest X-ray source is nof associated with MW(C 297 itself.," Using, we have resolved the X-ray emission from objects surrounding MWC 297 and found that the brightest X-ray source is not associated with MWC 297 itself."989 Furthermore. we have shown that the peak oft observed flaring is inconsistent with the position ADWC 297.," Furthermore, we have shown that the peak of the observed flaring is inconsistent with the position of MWC 297."990 Tustead. it is most likely due to a late-ty source iu the ITerbie Be stars field.," Instead, it is most likely due to a late-type source in the Herbig Be star's field."991 The study by Stelzer et al. (, The study by Stelzer et al. (9922003) oulate-tvpe stars. as well as our coronoeraphic study of tle vou18,"2003) onlate-type B stars, as well as our coronographic study of the young"9932003) oulate-tvpe stars. as well as our coronoeraphic study of tle vou18o,"2003) onlate-type B stars, as well as our coronographic study of the young"994 observed NCC 1313 and is X-ray emitting sources on 2002 October 13 (or 20592 sec with the ACIS detector (Carmireetal.2003)., observed NGC 1313 and its X-ray emitting sources on 2002 October 13 for 20592 sec with the ACIS detector \citep{Gar03}.995. The aimpoiut of the detector fell ou the 53 CCD., The aimpoint of the detector fell on the back-illuminated S3 CCD.996 We extracted a lel| curve of tie background. masking out all bright sources detectable by eve. to search [or aly sigatures of [kwing events known to affect. back-illumiuatecd CCDs (ChaucraProposers’(μίας2002).," We extracted a light curve of the background, masking out all bright sources detectable by eye, to search for any signatures of flaring events known to affect back-illuminated CCDs \citep{Chan02}."997. No strong lares were cetectecd but a few high points were eliminated. reducing the good exposure ime to 19902 sec.," No strong flares were detected but a few high points were eliminated, reducing the good exposure time to 19902 sec."998" We extracted the events at the loc:itiou of SN]LOTSIN. about 57.88 off-axis. in an aperture of radius 12"" which encloses 2985€ of the j»olnt sprea function (ChandraProposers!Guide2002)."," We extracted the events at the location of SN1978K, about $'$ .88 off-axis, in an aperture of radius $''$ which encloses $>$ of the point spread function \citep{Chan02}."999. The background. was obtained from an anuulus sur'ouudiug the source., The background was obtained from an annulus surrounding the source.1000 The net count rate was 20.1 10270.00:) counts 1, The net count rate was $\sim$ $\pm$ 0.003 counts $^{-1}$.1001 A response matrix was constructed specilic to tie olf-axis angle of 9NI978lIx. The matrix was corrected for the time-dependent absorption using t1e fitted functional form which depeuds upou a single parameter. the time from lauuch: for this observation. the time from launch was 1179 days (Plucinskyetal.2003).," A response matrix was constructed specific to the off-axis angle of SN1978K. The matrix was corrected for the time-dependent absorption using the fitted functional form which depends upon a single parameter, the time from launch; for this observation, the time from launch was 1179 days \citep{Plu03}."1002. Given the sharp point spread fuuction of the mirrors. we must be concerned with possible event pileup even though the large o[-axis angle 1jtigates the effects of pileup cousiderably.," Given the sharp point spread function of the mirrors, we must be concerned with possible event pileup even though the large off-axis angle mitigates the effects of pileup considerably."1003 The spectral analysis was undertaken using a pileup mode| (Davis200!L)., The spectral analysis was undertaken using a pileup model \citep{Davis01}.1004. The resulting fit iudicated low or zero pileup., The resulting fit indicated low or zero pileup.1005 Cuven the broader point sj»read Ltuictio rol the mirrors. there is no pileup of the EPIC-pu or MOS spectra from that instruiment.," Given the broader point spread function of the mirrors, there is no pileup of the EPIC-pn or MOS spectra from that instrument."1006 The resilts of the fits to tle auc spect‘a. to be described )elow. are ve'y similar. as one expects for observations separated by ~2 years of a slowly-evolving «):bject.," The results of the fits to the and spectra, to be described below, are very similar, as one expects for observations separated by $\sim$ 2 years of a slowly-evolving object."1007 The siuilarity sup»orts the conclusion of OW OL zero pileup., The similarity supports the conclusion of low or zero pileup.1008 We urther it the spectra witl both XSPEC (Arnaud1996) aud the Sherpa fitting eugiue (Freemal.Doe.Sielmlelnowska2001 and obtaiued identical results., We further fit the spectra with both XSPEC \citep{Arn96} and the Sherpa fitting engine \citep{FDS01} and obtained identical results.1009" We first usec the bes-fit moclels derived from tle ROSAT PSPC and SIS/CIS obse""vallons of SNLOTSIL w1ich consistecl of absorbect. sinele-component continuum models (R93. al. 199[.. S99)."," We first used the best-fit models derived from the PSPC and SIS/GIS observations of SN1978K which consisted of absorbed, single-component continuum models (R93, \citealt{Petre94}, S99)."1010 We lixec the model parameters a the previously-determinecd values. but it was üunmecdiately. evi(ent tha the models no loiger provided a good fit. instead. vieldiug a \7/v = D+) ," We fixed the model parameters at the previously-determined values, but it was immediately evident that the models no longer provided a good fit, instead yielding a ${\chi}^2/{\nu}$ $\geq$ 3."1011We then allowed the moclel parameters ο vary: tle resultiug moclel fits remained poor., We then allowed the model parameters to vary; the resulting model fits remained poor.1012 Figure shows the [it using a siugle absorbed BRavi1ond-Sultli mocel., Figure \ref{oldfit} shows the fit using a single absorbed Raymond-Smith model.1013 Clearly a sinele-component model is Inaclequate as it cau not simultaueous vlt theeulssion iu the 0.7-1.2 keV region as well as the apparent bard component at energies aJONre 2-3 keV. That theROSAT PSPC inodels do not provide a good fit is not especially surprIIT.g glvel the softer respouse of that detector., Clearly a single-component model is inadequate as it can not simultaneously fit the emission in the 0.7-1.2 keV region as well as the apparent hard component at energies above $\sim$ 2-3 keV. That the PSPC models do not provide a good fit is not especially surprising given the softer response of that detector.1014 That the spectra do not provide a good Li may be explained as either au iucrease in [lux above2, That the spectra do not provide a good fit may be explained as either an increase in flux above21015CCC coustraints. aud to understaud our ability to reconstruct the mass distribution from the galaxy data.,"CCC constraints, and to understand our ability to reconstruct the mass distribution from the galaxy data."1016 Iteparameterizine the observable galaxies by their (dimeusionless) aneular diameter clisalice d rather than redsift 2 makes it clear how he WL CCC method cau provide mocel-iudeperdeut geonetric colπίταus on ο., Reparameterizing the observable galaxies by their (dimensionless) angular diameter distance $d$ rather than redshift $z$ makes it clear how the WL CCC method can provide model-independent geometric constraints on $\Omega_k$.1017 Τιis methodoogy also reveals that CCC daa ls degenerate 1udderalterations to lub aud lud by quadratic [tuctious of d. if ο is free.," This methodology also reveals that CCC data is degenerate underalterations to $\ln b$ and $\ln d$ by quadratic functions of $d$, if $\Omega_k$ is free."1018" A)art from three exact degeneracies. hec""'Oss-COITeallo1 strength. )eing a joint function of the leus and source dista1Ces. is very efficient. at ooduciug cle'oupled estivates of O,. the distauces d; aud the bias factors Dd}. and cau even deterine shea' calibration facors f; lor each source plane wit1 modest degracdalous ol factor 352 in cosuological accWacy."," Apart from three exact degeneracies, the cross-correlation strength, being a joint function of the lens and source distances, is very efficient at producing decoupled estimates of $\Omega_k$, the distances $d_i$ and the bias factors $B_i$, and can even determine shear calibration factors $f_i$ for each source plane with modest degradations of factor $\lesssim 2$ in cosmological accuracy."1019 With sufficieutly accurate p0tometLic 'eshifts. the same survey data used for weak leislug may be used to determiue the ralisve‘se BAO scale.," With sufficiently accurate photometric redshifts, the same survey data used for weak lensing may be used to determine the transverse BAO scale."1020" This gives additional modelLincependent constraints coii d that are of lower orecislo 1ithan he CCC data. but are completely free of degene""cvy."," This gives additional model-independent constraints on $d$ that are of lower precision than the CCC data, but are completely free of degeneracy."1021" Such a comjued CCC-BAO strvey shokd s""jeld uicertaluties of e0.0L ος Q4."," Such a combined CCC-BAO survey should yield uncertainties of $\approx0.04 f_{\rm1022 sky}^{-1/2}$ on $\Omega_k$."1023 We 'elterate that such ¢Oustralits are competely idepeudeut of any assumptions about tle matte-energy content of je. Uuiverse. any biases Il photonetric redshifis. Or in [act any aleratious to the Friecdinauu equation sor the deecjon ecuations for light.," We reiterate that such constraints are completely independent of any assumptions about the matter-energy content of the Universe, any biases in photometric redshifts, or in fact any alterations to the Friedmann equations or the deflection equations for light."1024 Tvey merely require that the inetri€ be applicable to wu Universe., They merely require that the Robertson-Walker metric be applicable to our Universe.1025 Given ie lack of viable dark-energy heories. it seenus prudent to seek cosnological infonusion tha Ολα]is valid even if the acceleration is attributable to an alteraion of General Relativity ‘ather t lalla ϱ'eviously innoticed stress-eneoy contribution.," Given the lack of viable dark-energy theories, it seems prudent to seek cosmological information that remains valid even if the acceleration is attributable to an alteration of General Relativity rather than a previously unnoticed stress-energy contribution."1026 In the spirit of couservatislh. Wre shoulc| note that the RW metric may uot be suliciently accurate.," In the spirit of conservatism, we should note that the RW metric may not be sufficiently accurate."1027 TIe clummpiness of the IHter-elle‘oy disributiou uay invalklate our adoptior of the fillecd-beam augular diameter distance. and a ujore sophisticated treatment may be required2001).," The clumpiness of the matter-energy distribution may invalidate our adoption of the filled-beam angular diameter distance, and a more sophisticated treatment may be required."1028 It is not clear to wlal extell sina]|-scale cltumipiuess can influence the a»parelnt shear of galaxy-scale images as surveyed across the entire σεντς is quite a different regije from studies of srongvy leused quasars or superuovae for which tle Dyer-Boeder clistauces have been most carefully st«lied., It is not clear to what extent small-scale clumpiness can influence the apparent shear of galaxy-scale images as surveyed across the entire sky—this is quite a different regime from studies of strongly lensed quasars or supernovae for which the Dyer-Roeder distances have been most carefully studied.1029 The cu'vature measurenmieit is limited by tje ability to break the WL CCC degeneracy.i. by the accuracy and redshift spau ofthe BAO or Sv» Weastl‘ements of auetular-diameter distauces.," The curvature measurement is limited by the ability to break the WL CCC degeneracy, by the accuracy and redshift span of the BAO or SN measurements of angular-diameter distances."1030 Our baseline coistraint asstunes oto-z BAO infonation or 0<2<3 over the full sky., Our baseline constraint assumes $z$ BAO information for $0<z<3$ over the full sky.1031 A massive spectroscopic BAO survey woud offer substantial tprovement. but Type Ia supernova studies would uot oler stbstautial improverent over ftll-sky. oto-: BAO unless systematic uncertainties in SNIa pea& magnitudes coul be brought wel below 1! aud the redshift ranee extended beyond 2= 2.," A massive spectroscopic BAO survey would offer substantial improvement, but Type Ia supernova studies would not offer substantial improvement over full-sky $z$ BAO unless systematic uncertainties in SNIa peak magnitudes could be brought well below 1% and the redshift range extended beyond $z=2$ ."1032 Incliion of lensing aud acoustic-scale informaion [rom the epochs of recombination or reionization would help significantly., Inclusion of lensing and acoustic-scale information from the epochs of recombination or reionization would help significantly.1033For the pulsation search the task of the analysis software was used with the x? test statistics.,For the pulsation search the task of the analysis software was used with the $\chi^{2}$ test statistics.1034 We used 10 phase bins for the x? test with a step size of 0.5 in units of Fourier resolution that lead us to a number of 5000 independent trials., We used 10 phase bins for the $\chi^{2}$ test with a step size of 0.5 in units of Fourier resolution that lead us to a number of 5000 independent trials.1035 By checking a frequency range from 3.76885824639721 to 3.76929999373175 Hz and using f——7.78x107? Hz s! we found the most probable frequency to be 3.76908389(3) Hz at a test statistic of x?=328.0 with 9 degree of freedom (see Figure 11))., By checking a frequency range from 3.76885824639721 to 3.76929999373175 Hz and using $\dot{f}=-7.78\times10^{-13}$ Hz $^{-1}$ we found the most probable frequency to be 3.76908389(3) Hz at a test statistic of $\chi^{2}=328.0$ with 9 degree of freedom (see Figure \ref{pow_spec}) ).1036 The number in the parenthesis is the uncertainty of the last digit of the quoted frequency which corresponds to the Fourier resolution., The number in the parenthesis is the uncertainty of the last digit of the quoted frequency which corresponds to the Fourier resolution.1037" We have also performed the analysis by combining Z2 test and H —test, where n is the numbers of harmonics (Buccheri et al."," We have also performed the analysis by combining $Z^{2}_{n}$ test and $H-$ test, where $n$ is the numbers of harmonics (Buccheri et al."1038" 1983; De Jager, Swanepoel, Raubenheimer 1989)."," 1983; De Jager, Swanepoel, Raubenheimer 1989)."1039" With this independent analysis, we reported the periodic signal with a frequency equal to the aforementioned."," With this independent analysis, we reported the periodic signal with a frequency equal to the aforementioned."1040" Using the H-test, we found that H is maximized for 2 harmonics."," Using the $H$ -test, we found that $H$ is maximized for 2 harmonics."1041" The calculated Z2 is 309.7 which implies a nominal chance probability of 4x10-92,", The calculated $Z^{2}_{2}$ is 309.7 which implies a nominal chance probability of $4\times10^{-62}$.1042 We have also repeated the analysis with different positions within the 9596 y—ray error circle adopted for barycentric correction., We have also repeated the analysis with different positions within the $95\%$ $\gamma-$ ray error circle adopted for barycentric correction.1043" Among all the tested positions, the X-ray position of rresults in the best test statistics."," Among all the tested positions, the X-ray position of results in the best test statistics."1044" For example, the nominal -—ray position provided by Abdo et al. ("," For example, the nominal $\gamma-$ ray position provided by Abdo et al. ("10452009a) results in a X? of 280.9 which is lower than that resulted from adopting the X-ray position.,2009a) results in a $\chi^{2}$ of 280.9 which is lower than that resulted from adopting the X-ray position.1046 This provides another support for the possible association between aandJ, This provides another support for the possible association between and.10472021+4026.. The y—ray pulse profile folded at the aforementioned period is shown in Figure 12., The $\gamma-$ ray pulse profile folded at the aforementioned period is shown in Figure \ref{lat_pulse}.1048" We found that hhas a double peaked light curve with the peak separation of 162° or 198°, depending on which peak is leading."," We found that has a double peaked light curve with the peak separation of $162^{\circ}$ or $198^{\circ}$, depending on which peak is leading."1049" In computing the pulsed fraction of the resultant light curve, we found that about ~ of the collected photons are pulsed."," In computing the pulsed fraction of the resultant light curve, we found that about $\sim$ of the collected photons are pulsed."1050 This light curve provides us a crucial input for modeling the emission geometry of this pulsar (see below)., This light curve provides us a crucial input for modeling the emission geometry of this pulsar (see below).1051'The reduction procedure followed is identical to that described in Vogt&Dopita(2010a) for their observations of SNR 1E 0102.2-7219.,The reduction procedure followed is identical to that described in \cite{Vogt10a} for their observations of SNR 1E 0102.2-7219.1052 The six fields were each reduced separately using the WiFeS reduction pipeline based upon IRAF! scripts., The six fields were each reduced separately using the WiFeS reduction pipeline based upon IRAF scripts.1053 This procedure has been described in Dopitaetal.(2010) and will not be repeated here., This procedure has been described in \citet[][]{dopita2010} and will not be repeated here.1054 The same calibration files were used for all fields., The same calibration files were used for all fields.1055 The frames used for flat fielding and tracing the slit positions consist of a dome wire and a sky wire (obtained using the coronographic science aperture with a 1.0 arc sec., The frames used for flat fielding and tracing the slit positions consist of a dome wire and a sky wire (obtained using the coronographic science aperture with a 1.0 arc sec.1056" wire stretched across the middle), a dome flat and sky flats (obtained using the science aperture."," wire stretched across the middle), a dome flat and sky flats (obtained using the science aperture."1057 All sky flats were taken at dusk of the first night., All sky flats were taken at dusk of the first night.1058 Bias and Ne Ar arc exposures for wavelength calibration were taken directly before or directly after the group of three exposures on each field., Bias and Ne Ar arc exposures for wavelength calibration were taken directly before or directly after the group of three exposures on each field.1059" Each of the three images of each of the six fields are bias subtracted, flat fielded and sky subtracted, atmospheric dispersion corrected, un-binned, and finally combined and reduced to a wavelength calibrated 3D data cube from which the spectra of individual regions can be extracted, or from which monochromatic images of the whole field can be formed."," Each of the three images of each of the six fields are bias subtracted, flat fielded and sky subtracted, atmospheric dispersion corrected, un-binned, and finally combined and reduced to a wavelength calibrated 3D data cube from which the spectra of individual regions can be extracted, or from which monochromatic images of the whole field can be formed."1060" We have not calibrated the data to absolute flux units using observations of photometric standard stars, as this is not required for the purposes of our dynamical study."," We have not calibrated the data to absolute flux units using observations of photometric standard stars, as this is not required for the purposes of our dynamical study."1061 'The issue of the bias subtraction deserves a special mention., The issue of the bias subtraction deserves a special mention.1062" The WiFeS bias frames are subject to small temporal variations, but have little or no fixed pattern noise."," The WiFeS bias frames are subject to small temporal variations, but have little or no fixed pattern noise."1063 Since each chip is read out in all four, Since each chip is read out in all four1064(2002).,(2002).1065 This is troublesome. from the observational point of view. because it implies that the choice of energv bands (which is usually done withoutpriori knowledge about the overall SED) max affect what we see observationally.," This is troublesome, from the observational point of view, because it implies that the choice of energy bands (which is usually done without knowledge about the overall SED) may affect what we see observationally."1066 To further complicate (he situation. itis nearly impossible. observationallv. to cleanly isolate a flare from the hierarchical structure.," To further complicate the situation, it is nearly impossible, observationally, to cleanly isolate a flare from the hierarchical structure."1067 We wish to thank Mark Ertmer and Ixie Li for assistance in data reduction. Feng Yuan and Markus Bottitcher for useful discussions and comments on the manuscript.," We wish to thank Mark Ertmer and Kie Li for assistance in data reduction, Feng Yuan and Markus Bötttcher for useful discussions and comments on the manuscript."1068 This research has made use of data obtained through5 the Hieh5 Enerey5. Astrophysics Science Archive Research Center Online Service. provided by the NASA/Goddard Space Flight Center.," This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center."1069 This work was supported in part bv the NASA grant. NAG5-13736., This work was supported in part by the NASA grant NAG5-13736.1070part of line of clusters of ealaxies along a filament.,part of line of clusters of galaxies along a filament.1071"components (D.Dp. D,,)). 3 is the electric current. and wis the circulation which has components (n.np.i65,).","components $\Br,\Bt,\Bp$ )), $\bj$ is the electric current, and $\bu$ is the circulation which has components $(\ur,\ut,\up)$."1072 In order to derive these equations. the density. p ds assumed to be constant (py= lg 3) everywhere. as is the kinematic viscosity 7 and the magnetic diffusivity η.," In order to derive these equations, the density $\rho $ is assumed to be constant $\rho_0 =1$ g $^{-3} $ ) everywhere, as is the kinematic viscosity $\nu$ and the magnetic diffusivity $\eta$."1073 This approximation is mace to simplify the numerical procedure slightlv. but can casily be removed.," This approximation is made to simplify the numerical procedure slightly, but can easily be removed."1074 Since. as it will be shown later. the tvpical values of the viscosity and magnetic cdilfusivitv used in the simulations are several orders of magnitude larger than in the sun. it seenis pointless to try and represent accurately the variation of these quantities in this first analysis.," Since, as it will be shown later, the typical values of the viscosity and magnetic diffusivity used in the simulations are several orders of magnitude larger than in the sun, it seems pointless to try and represent accurately the variation of these quantities in this first analysis."1075 Using the following new svstem of units: where By is the typical strength of the radial field in the interior. the equations become: with where all the quantities are now dimensionless ancl e. is the unit vector. parallel to the rotation axis.," Using the following new system of units: where $B_0$ is the typical strength of the radial field in the interior, the equations become: with where all the quantities are now dimensionless and $\ez$ is the unit vector parallel to the rotation axis."1076" The Ekman numbers £, and £,, represent the ratios of the rotation timescale to the cilfusive timescales. ancl the Elsasser nuniber Ais the ratio of the typical amplitude of the Lorentz force to that of the Coriolis force."," The Ekman numbers $\Enu$ and $\Eeta$ represent the ratios of the rotation timescale to the diffusive timescales, and the Elsasser number $\Lambda$ is the ratio of the typical amplitude of the Lorentz force to that of the Coriolis force."1077" Ehe system is solved for (nj.no.ns) and CD.Do.D,,) by solving the azimuthal component of the momentum equation. the azimuthal component of the vorticity equation. the integrated induction. equation and. the azimuthal component of the induction equation. together with the mass conservation equation and the solenoidal condition: those equations are where w=Vw is the vorticity."," The system is solved for $(\ur,\ut,\up)$ and $(\Br,\Bt,\Bp)$ by solving the azimuthal component of the momentum equation, the azimuthal component of the vorticity equation, the integrated induction equation and the azimuthal component of the induction equation, together with the mass conservation equation and the solenoidal condition; those equations are where $\bomega = \curl \bu$ is the vorticity."1078 1n this first analysis. the boundary conditions chosen for the system are the simplest possible ones that still guarantee the existence of a solution: as a result. they are not necessarily the most accurate representation of the dynamics of the sun (and in particular of the interface of the tachocline with the convection zone).," In this first analysis, the boundary conditions chosen for the system are the simplest possible ones that still guarantee the existence of a solution; as a result, they are not necessarily the most accurate representation of the dynamics of the sun (and in particular of the interface of the tachocline with the convection zone)."1079 The ellects of the boundary conditions on the system. ancl possible improvements. are discussed in Section 6.3 The boundary conditions are summarized in Fig. 1:," The effects of the boundary conditions on the system, and possible improvements, are discussed in Section \ref{sec:discmod}1080 The boundary conditions are summarized in Fig. \ref{fig:bc};"1081 the MIID equations presented in equations (5)) are solved in a region located. between two spherical boundaries. impermeable ancl no-slip boundaries so that the radial aud the latituclinal components of the circulation vanish on the boundaries. and the azimuthal component of the circulation is given by the rotation of the boundaries.," the MHD equations presented in equations \ref{eq:eqsmhd}) ) are solved in a region located between two spherical boundaries, impermeable and no-slip boundaries so that the radial and the latitudinal components of the circulation vanish on the boundaries, and the azimuthal component of the circulation is given by the rotation of the boundaries."1082" As a result. the angular-velocitv. perturbation is O-CO04© on the inner boundary. where Qi, is an eigenvalue of the problem. (see below) and Q=OG,(80)QO. on the outer boundary. where and Qo,= LOTO. a2=a,0.15 (according to the observations presentec by Schou ct al. ("," As a result, the angular-velocity perturbation is $\tilde{\Omega} =1083 \Omega_{\rm in} - \Omega_{\rm c}$ on the inner boundary where $\Omega_{\rm in}$ is an eigenvalue of the problem (see below) and $\tilde{\Omega} = \Omega_{\rm cz}(\theta) - \Omega_{\rm c}$ on the outer boundary, where and $\Omega_{\rm eq} = 1.07 \Oc$ , $a_2 = a_4 = 0.15$ (according to the observations presented by Schou et al. ("10841998)).,1998)).1085" On the equatorial plane. svmimetry arguments determine the behaviour of the solutions: sj—0. 2.=0. 00/08—0 and I,,=0."," On the equatorial plane, symmetry arguments determine the behaviour of the solutions: $\ut = 0$, $\Br = 0$, $\ptl \tilde{\Omega}/\ptl \theta = 0 $ and $\Bp = 0$."1086" On the poles. regularity conditions impose ""0= 0. By=0. 00/00=0 and D,=0."," On the poles, regularity conditions impose $\ut = 0$ , $\Bt = 0$, $\ptl \tilde{\Omega}/\ptl \theta = 0$ and $\Bp = 0$."1087 Phe bounding spheres are assumed to be imperfectIy conducting and the region outside Juda] Supports no fluid motion. which implies that the licld in that region satislies V7.B.—0. with conditions that DBs Oat infinity. anc that the field structure becomes purely dipolar as roo»0.," The bounding spheres are assumed to be imperfectly conducting and the region outside $[\rin,\rout]$ supports no fluid motion, which implies that the field in that region satisfies $\grad^2 \bB =0$, with conditions that $\bB \rightarrow 0$ at infinity, and that the field structure becomes purely dipolar as $r\rightarrow 0$."1088 The amplitude of the racial component of the field near the poles is fixed. such that Bin.0—0)=By.," The amplitude of the radial component of the field near the poles is fixed such that $B_r(\rin,\theta=0) =1089 B_0$."1090" This uniquely determines the boundary conditions for the magnetic field. as a relation between D, and D, on the one hand. and between ο, anc OD,für on the other hand."," This uniquely determines the boundary conditions for the magnetic field, as a relation between $\Br$ and $\Bt$ on the one hand, and between $\Bp$ and $\ptl \Bp/\ptl r$ on the other hand."1091 As a result. of these conditions. when he svstem is not rotating. the solution for the magnetic 101 is a purely poloidal dipolar structure with a field streneth varving as 1/677.," As a result of these conditions, when the system is not rotating, the solution for the magnetic field is a purely poloidal dipolar structure with a field strength varying as $1/r^3$."1092 When the svstem is rotating. hese boundary conditions drive a meridional How. through Ekman-llartmann pumping.," When the system is rotating, these boundary conditions drive a meridional flow through Ekman-Hartmann pumping."1093 This Low is used to mimic the xwoclinicl« driven How predicted to occur in the tachocline (Spiegel Zahn. 1992. Gough AMelnotvre. 1998). in order to study its interaction with he imposed. magnetic field.," This flow is used to mimic the baroclinicly driven flow predicted to occur in the tachocline (Spiegel Zahn, 1992, Gough McIntyre, 1998), in order to study its interaction with the imposed magnetic field."1094" Phe principal caveat ofthis approach is that the typical velocities ofthe meridional Low scales with £, and {κ as an Ekman-LHirtmaan Iow rather", The principal caveat ofthis approach is that the typical velocities ofthe meridional flow scales with $\Enu$ and $\Eeta$ as an Ekman-Hartmann flow rather1095 Ho-ouitting," $\sf{\alpha}$ \citep{Schmitt03,Veilleux03,Whittle04,Ramos06}. \citep{Silk98,Kauffmann03,Hopkins06,Schawinski07,Schawinski09}."1096 Oue of the inost promising wavs to study these interactions is through soft N-ravs(Crenshawetal.1999.2003:Daietal. 2008).," One of the most promising ways to study these interactions is through soft X-rays \citep{Crenshaw99,Crenshaw03,Dai08}. ."1097. In the unified picture of AGN (Antonucci-1993) the soft N-rav spectra of Tvpe-2 Seyfert galaxics are expected tobe affected by. emission and scattering from the mecdimm. which is also stronely influenced by the unclear continu.," In the unified picture of AGN \citep{Antonucci93} the soft X-ray spectra of Type-2 Seyfert galaxies are expected to be affected by emission and scattering from the medium, which is also strongly influenced by the nuclear continuum."1098 However. soft N-ravs can also be produced through mechanical heating. as in shocks dviven by supernova explosions in nuclear star-forming regious.," However, soft X-rays can also be produced through mechanical heating, as in shocks driven by supernova explosions in nuclear star-forming regions."1099 Indeed. it is quite plausible that both effects are important ΑΛ: Sakoetal. 2000)).," Indeed, it is quite plausible that both effects are important 3; \citealt{Sako00}) )."1100 This soft X-ray cuaission{ ]provides the opportunity to obtain an XN-rav diagnostic for the plivsical properties of the interacting interstellar edit [ISM| (YoungetetTheal.2006:ealaxyEvansctIraenmier 200," This soft X-ray emission provides the opportunity to obtain an X-ray diagnostic for the physical properties of the interacting interstellar medium [ISM] \citep{Young01,Yang01,Ogle00,Ogle03,Bianchi06,Evans06,Kraemer08}."11018).. Ablyk 573 hasheen extensively studied bv nuu authors., The galaxy Mrk 573 has been extensively studied by many authors.1102 Its active nucleus is hosted iuan (R)JSAD(rs)0] ΠΟΤΟ is well-known for its extended. richly structured circumuutuclear eission-Iue," Its active nucleus is hosted inan (R)SAB(rs)0+ 573 is well-known for its extended, richly structured circumnuclear emission-line"11030) non-Gaussian contributions. the probability of overcoming the collapse threshold becomes higher.,") non-Gaussian contributions, the probability of overcoming the collapse threshold becomes higher."1104 As a consequence. the formation of high-mass haloes is enhanced and anticipated when {xcc0.," As a consequence, the formation of high-mass haloes is enhanced and anticipated when $f_{\rm NL}>0$."1105 Fig., Fig.1106 3 shows that the mass function can be increased by a factorof 10 at 2=13 for haloes with mass AZ~1071A7. when compared to the standard scenario.," \ref{fig:1a} shows that the mass function can be increased by a factor of 10 at $z=13$ for haloes with mass $M\sim110710^{11}M_{\odot}$ when compared to the standard scenario."1108" We should however remark that high-mass haloes (AJ>10A7. ) at early cosmological epochs are rare events. as shown also by the small number density at >=19 in the reference case. n(2»10AL.)25.10 ""IMpe*."," We should however remark that high-mass haloes $M>10^{9}M_{\odot}$ ) at early cosmological epochs are rare events, as shown also by the small number density at $z=13$ in the reference case, $n(>10^{9}M_{\odot})\lesssim11095 \times 10^{-3}$ $^{3}$."1110 Then this effect is expected to have a little impact on integrated quantities as the total ionized fraction ot the IGM optical depth., Then this effect is expected to have a little impact on integrated quantities as the total ionized fraction ot the IGM optical depth.1111 Unlike the local model. the scale dependence of non-Gaussianity increases the abundance of the low-mass haloes by a factor of ~10 at 2=13 when compared to the standard case.," Unlike the local model, the scale dependence of non-Gaussianity increases the abundance of the low-mass haloes by a factor of $\sim 10$ at $z=13$ when compared to the standard case."1112 The opposite applies for (νι<0., The opposite applies for $f_{\rm NL}<0$.1113 As already noticed by Matarreseetal.(2000). [see also Verdeetal.(2001):Grossi(2007 )]]. this effect is more evident at early cosmological epochs. exactly when the process of IGM ionization starts.," As already noticed by \cite{matarrese2000} [see also \cite{verde2001,grossi2007}] ], this effect is more evident at early cosmological epochs, exactly when the process of IGM ionization starts."1114 For this reason a non-Gaussian distribution of the primordial density field can affect the way in which reionization occurs. leaving its imprints on it. as we will investigate in the next sections.," For this reason a non-Gaussian distribution of the primordial density field can affect the way in which reionization occurs, leaving its imprints on it, as we will investigate in the next sections."1115 In this section. we briefly review the main assumptions underlying the analytic model adopted to describe the process of cosmic reionization.," In this section, we briefly review the main assumptions underlying the analytic model adopted to describe the process of cosmic reionization."1116 This model is based on the approach proposed by Avelino&Liddle(2006) [see also Haiman&Holder(2003):Chenetal.(2003) for further details]: our implementation. however. differs in some aspects which will be discussed later.," This model is based on the approach proposed by \cite{avelino2006} [see also \cite{haiman2003,chen2003} for further details]; our implementation, however, differs in some aspects which will be discussed later."1117 In this model. the statistical properties of the ionized regions are related to the hierarchical growth of the ionizing sources through simple assumptions on how the galaxies ionize the IGM and on how the IGM recombines.," In this model, the statistical properties of the ionized regions are related to the hierarchical growth of the ionizing sources through simple assumptions on how the galaxies ionize the IGM and on how the IGM recombines."1118 A one-to-one correspondence between the distribution of galaxies and ΠΠ regions is established. such that a single galaxy of mass Ad... can ionize a region of mass Muu=CM.," A one-to-one correspondence between the distribution of galaxies and HII regions is established, such that a single galaxy of mass $M_{\rm gal}$ can ionize a region of mass $M_{\rm HI\!I}=\zeta1119M_{\rm gal}$."1120 Here ¢ represents the ionization efficiency of the galaxy. and it is strictly dependent on the nature of the ionizing sources.," Here $\zeta$ represents the ionization efficiency of the galaxy, and it is strictly dependent on the nature of the ionizing sources."1121 We will take it as a constant. fixed in such a way that reionization ends at 2= 6.5.," We will take it as a constant, fixed in such a way that reionization ends at $z=6.5$ ."1122 Since at high + the cooling of the gas becomes efficient in aloes having a virial temperature Zz:101 K. unlike in Avelino&Liddle (2006).. in our analysis we consider only type Ia (10! K Tox9510! Kyand type Ib CE>9107 Ky haloes. neglecting he contribution of the type II sources. which would correspond to aaloes with 400 K T107 K. We recall that the distinction between the halo types is related to the way in which they impact he IGM: type Ia sources can grow only in neutral regions. while ype Ib haloes can appear also in ionized regions.," Since at high $z$ the cooling of the gas becomes efficient in haloes having a virial temperature $T\ge 10^4$ K, unlike in \cite{avelino2006}, in our analysis we consider only type Ia $10^4$ K $\le T \le 9 \times10^4$ K) and type Ib $T > 9\times 10^4$ K) haloes, neglecting the contribution of the type II sources, which would correspond to haloes with $400$ K $\le T \le 10^4$ K. We recall that the distinction between the halo types is related to the way in which they impact the IGM: type Ia sources can grow only in neutral regions, while type Ib haloes can appear also in ionized regions."1123 Consequently hey affect differently the ionization phases of IGM., Consequently they affect differently the ionization phases of IGM.1124" The total collapsed fraction πως) at different redshifts can be computed using eq.(6)): where po is the present-day matter density and μμ(>) is the minimum mass corresponding to the virial temperature 7. which can be computed by inverting the relation proposed by namely: In the previous equation. A. represents the virial overdensity at redshift 2 and £25, is the matter density parameter at redshift 2."," The total collapsed fraction $F_{\rm coll}(z)$ at different redshifts can be computed using \ref{eq:1c}) ): where $\bar{\rho}_{0}$ is the present-day matter density and $M_{\rm1125 min}(z)$ is the minimum mass corresponding to the virial temperature $T$, which can be computed by inverting the relation proposed by \cite{barkana2001}, namely: In the previous equation, $\Delta_{\rm c}$ represents the virial overdensity at redshift $z$ and $\Omega_{\rm m}^{z}$ is the matter density parameter at redshift $z$."1126" Consequently. the collapsed fractions in Ta and Ib haloes are given by where Adin, and Adin. are the minimum masses for Ib and Tu sources. obtained using in eq.(10)) Z=9«107 and 107 K. respectively."," Consequently, the collapsed fractions in Ia and Ib haloes are given by where $M_{{\rm min},Ib}$ and $M_{{\rm min},Ia}$ are the minimum masses for Ib and Ia sources, obtained using in \ref{eq:3a}) ) $T= 9\times112710^4$ and $10^4$ K, respectively."1128 The action of the ionizing sources is smoothed down by the recombination of the IGM. here considered as a homogeneous gs.," The action of the ionizing sources is smoothed down by the recombination of the IGM, here considered as a homogeneous gas."1129 The recombination rate is linearly dependent on the IGM clumping factor Cyy=<ninfoconumfor which. following Haiman&Bryan (2006)... we assume a redshift evolution modeled as: being 2 a free parameter.," The recombination rate is linearly dependent on the IGM clumping factor $C_{\rm HI\!I}=<n_{\rm HI\!I}^2>/<n_{\rm HI\!I}>^2$, for which, following \cite{haiman2006}, we assume a redshift evolution modeled as: being $\beta$ a free parameter."1130" As shown by Avelino&Liddle (2006)... the predicted reionization history of the universe has Significant uncertainties introduced by the poor knowledge of the z-dependence of the clumping factor, which cannot be robustly constrained even considering the 3-year WMAP results for the reionization optical depth."," As shown by \cite{avelino2006}, the predicted reionization history of the universe has significant uncertainties introduced by the poor knowledge of the $z$ -dependence of the clumping factor, which cannot be robustly constrained even considering the 3-year WMAP results for the reionization optical depth."1131 Since they found good consistency between predicted and observed optical depths irrespectively of the amount of primordial non-Gaussianity in the models they consider. we decide to set ή=0.," Since they found good consistency between predicted and observed optical depths irrespectively of the amount of primordial non-Gaussianity in the models they consider, we decide to set $\beta=0$."1132 In this case. the z-dependence of C'Hu is neglected and 6Ἡν= 10.," In this case, the $z$ -dependence of $C_{\rm HI\!I}$ is neglected and $C_{\rm HI\!I}=10$ ."1133 The impact of the assumption of a constant clumping factor will be discussed later., The impact of the assumption of a constant clumping factor will be discussed later.1134" The probability that a photon emitted at a given cosmological epoch σι(1) is still ionizing at 2<z; can be written as with /,=pCanai). being aj; the recombination coefficient of HI (=2.610£7 em""/s at Z—10! K) and ngi(z)=LSSQuoAT/0.022(1|z) em? the hydrogen density at redshift +."," The probability that a photon emitted at a given cosmological epoch $z_{\rm i}(t_{\rm i})$ is still ionizing at $z<z_{\rm i}$ can be written as with $t_{\rm r}=\alpha_{\rm B} C_{\rm HI\!I} n_{\rm HI}(t_{\rm i})t_{\rm i}^{2}$, being $\alpha_{\rm B}$ the recombination coefficient of HI $=2.6 \times113510^{-13}$ $^{3}$ /s at $T=10^4$ K) and $n_{\rm HI}(z)=1.88\Omega_{\rm b0}h^{2}/0.022(1+z)^{3}$ $^3$ the hydrogen density at redshift $z$."1136 Then. the filling factor Hus) at a given cosmological epoch is where the ionizing efficiency is assumed to be the same for the different types of haloes.," Then, the filling factor $F_{\rm HI\!I}(z)$ at a given cosmological epoch is where the ionizing efficiency is assumed to be the same for the different types of haloes."1137 We notice that the different nature of type Ib and Ta sources appears in the right-hand sideof eq.¢1+)). where the (μι) factor explicitly considers that typeIu haloes form only in neutral regions.," We notice that the different nature of type Ib and Ia sources appears in the right-hand sideof \ref{eq:2e}) ), where the $(1-F_{\rm HI\!I})$ factor explicitly considers that typeIa haloes form only in neutral regions."1138"a gradient can be preserved or even enhanced through a merger. this is only true when they consider a “primary” galaxy merging with a ""companion"" of much steeper slope.","a gradient can be preserved or even enhanced through a merger, this is only true when they consider a “primary” galaxy merging with a “companion” of much steeper slope."1139 In the realistic case of a merger between two galaxies of equal initial slope. they find the remnant to have 0.6 times this initial slope.," In the realistic case of a merger between two galaxies of equal initial slope, they find the remnant to have 0.6 times this initial slope."1140 Repeated mayor mergers would thus very quickly wash out the initial gradients., Repeated major mergers would thus very quickly wash out the initial gradients.1141 If the UV colour gradients are driven by helium abundance. the absence of a regeneration mechanism means that they provide a stronger constraint on the importance of dry mergers in the formation of ellipticals than do the metallicity gradients.," If the UV colour gradients are driven by helium abundance, the absence of a regeneration mechanism means that they provide a stronger constraint on the importance of dry mergers in the formation of ellipticals than do the metallicity gradients."1142 The weak AGN and Liners shows very little difference in their UV properties from the main sample. particularly once the central 3 aresec aperture is taken out.," The weak AGN and Liners shows very little difference in their UV properties from the main sample, particularly once the central 3 arcsec aperture is taken out."1143 It appears that a weak AGN has little effect on the global UV properties., It appears that a weak AGN has little effect on the global UV properties.1144 The SO/SAO/SABO galaxies differ from the main sample in their NUV properties. in particular in Vy;. which is strongly positive in many of this sample.," The S0/SA0/SAB0 galaxies differ from the main sample in their NUV properties, in particular in $\nabla_{NJ}$, which is strongly positive in many of this sample."1145 However the FUV properties of this sample follow largely the same correlations as the main sample. and it seems that the FUV excess is unrelated to any ongoing star formation activity.," However the FUV properties of this sample follow largely the same correlations as the main sample, and it seems that the FUV excess is unrelated to any ongoing star formation activity."1146 In agreement with previous authors we find that early-type galaxies show an excess of flux in the GALEX far ultraviolet band. compared with that expected from a normal old stellar population.," In agreement with previous authors we find that early-type galaxies show an excess of flux in the GALEX far ultraviolet band, compared with that expected from a normal old stellar population."1147 There is a considerable variation in the magnitude of this excess. and a complex set of dependencies upon the dynamical and stellar population parameters of the galaxy.," There is a considerable variation in the magnitude of this excess, and a complex set of dependencies upon the dynamical and stellar population parameters of the galaxy."1148 The excess shows a stronger correlation with the central velocity dispersion than with the absolute magnitude of the galaxy., The excess shows a stronger correlation with the central velocity dispersion than with the absolute magnitude of the galaxy.1149 It also shows a definite correlation with [a/Fe] in the old stellar population. and a weaker correlation with [Z/H].," It also shows a definite correlation with $\alpha$ /Fe] in the old stellar population, and a weaker correlation with [Z/H]."1150 The FUV excess is more centrally concentrated than. the underlying old stellar population in all galaxies except for currently starforming systems and recent major merger remnants., The FUV excess is more centrally concentrated than the underlying old stellar population in all galaxies except for currently starforming systems and recent major merger remnants.1151 The logarithmic (FUV-NUV) colour gradient Vpv is overwhelmingly positive. is typically 0.6 magnitudes/dex. and shows a weak positive correlation with the central internal velocity dispersion oy. but it correlates much more weakly with the strength of the FUV excess itself.," The logarithmic (FUV-NUV) colour gradient $\nabla_{FN}$ is overwhelmingly positive, is typically 0.6 magnitudes/dex, and shows a weak positive correlation with the central internal velocity dispersion $\sigma_0$, but it correlates much more weakly with the strength of the FUV excess itself."1152 It does not correlate significantly with any stellar population parameter., It does not correlate significantly with any stellar population parameter.1153 Most galaxies show a negative colour gradient in (NUV-J). which can be attributed to a metallicity gradient.," Most galaxies show a negative colour gradient in (NUV-J), which can be attributed to a metallicity gradient."1154" This gradient is strongest in galaxies with loge,<2.35 although we do not sample galaxies with logay<2.1.", This gradient is strongest in galaxies with $\log{\sigma_0} < 2.35$ although we do not sample galaxies with $\log{\sigma_0} < 2.1$.1155" The relationship between the Sérrsic index η, and [o/Fe] in the stellar population found by Marino et al. "," The relationship between the Sérrsic index $n$, and $\alpha$ /Fe] in the stellar population found by Marino et al. ("11562011) appears to be driven largely by the NUV band. and appears to be related to the strong UV to IR colour gradients in galaxies with low c. and the correlation between συ and [o/Fe].,"2011) appears to be driven largely by the NUV band, and appears to be related to the strong UV to IR colour gradients in galaxies with low $\sigma_0$, and the correlation between $\sigma_0$ and $\alpha$ /Fe]."1157 When we treat the FUV excess as a separate component. we find that it is well fit with a Sérrsic model profile with the Sérrsic index » in the range 0 to 3.5.," When we treat the FUV excess as a separate component, we find that it is well fit with a Sérrsic model profile with the Sérrsic index $n$ in the range 0 to 3.5."1158 This component can contribute up to of the total ΕΙΝ flux rom the galaxy., This component can contribute up to of the total FUV flux from the galaxy.1159 The most likely origin of the FUV excess is in a population of old. probably metal poor. stars which coexist with he optically dominant old metal-rich population in the centres o ellipticals.," The most likely origin of the FUV excess is in a population of old, probably metal poor, stars which coexist with the optically dominant old metal-rich population in the centres of ellipticals."1160 Helium abundance is a plausible candidate for the property of the stellar population which drives both the variation in the strength of the excess. and the comparatively uniform radia gradients in that strength.," Helium abundance is a plausible candidate for the property of the stellar population which drives both the variation in the strength of the excess, and the comparatively uniform radial gradients in that strength."1161 Helium sedimentation in a cooling plasma in the very early stages of galaxy formation is a possible mechanism to establish such a gradient. although the timescale required to set up such a gradient by this process may be unfeasibly long.," Helium sedimentation in a cooling plasma in the very early stages of galaxy formation is a possible mechanism to establish such a gradient, although the timescale required to set up such a gradient by this process may be unfeasibly long."1162" The persistence of strong gradients to the present day provides strong limits of the importance of major ""dry"" mergers in the formation of ellipticals.", The persistence of strong gradients to the present day provides strong limits of the importance of major “dry” mergers in the formation of ellipticals.1163 GALEX is a NASA Small Explorer., GALEX is a NASA Small Explorer.1164 We acknowledge support rom NASA for construction. operation. and science analysis for he GALEX mission. developed in cooperation with the Centre ational d'Etudes Spatiales of France and the Korean Ministry of Science and Technology.," We acknowledge support from NASA for construction, operation, and science analysis for the GALEX mission, developed in cooperation with the Centre National d'Etudes Spatiales of France and the Korean Ministry of Science and Technology."1165 This publication makes use of data ooduets. from. the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology. unded by the National Aeronautics and Space Administration and the National Science Foundation.," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."1166 This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated bv the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."1167 DC and AMK acknowledge support from the Science and Technology Facilities Council (STFC) under grant ST/H/002391/1., DC and AMK acknowledge support from the Science and Technology Facilities Council (STFC) under grant ST/H/002391/1.1168 We acknowledge support for JK in this research from a Nuffield Foundation Science Bursary., We acknowledge support for JK in this research from a Nuffield Foundation Science Bursary.1169 We thank Dr. Sue Percival and Dr. Phil James for many helpful discussions and comments throughout the course of this project. and the referee for a careful reading and suggestions which have greatly improved the paper.," We thank Dr. Sue Percival and Dr. Phil James for many helpful discussions and comments throughout the course of this project, and the referee for a careful reading and suggestions which have greatly improved the paper."1170"The light curves show a persistent level of emission without significant variations of intensity, except for the eclipse periods, which are clearly seen in all the observations.","The light curves show a persistent level of emission without significant variations of intensity, except for the eclipse periods, which are clearly seen in all the observations."1171 The persistent count rate decreased steadily by from the first to the fourth observation., The persistent count rate decreased steadily by from the first to the fourth observation.1172" All the observations show a soft spectral component between 0.3 and —3 keV, which we fitted with a NS atmosphere model."," All the observations show a soft spectral component between 0.3 and $\sim$ 3 keV, which we fitted with a NS atmosphere model."1173 The bolometric flux of the thermal component decreased by between the first and the last observation., The bolometric flux of the thermal component decreased by between the first and the last observation.1174" In parallel, we observe a decrease in the effective surface temperature of the NS from 120 to 109 eV, which we interpret as the cooling of the NS crust towards thermal equilibrium with the core, after having been heated by accretion during the 24 year outburst."," In parallel, we observe a decrease in the effective surface temperature of the NS from 120 to 109 eV, which we interpret as the cooling of the NS crust towards thermal equilibrium with the core, after having been heated by accretion during the 24 year outburst."1175" We detected only a decrease in flux and a decrease in temperature between the last two observations, indicating that the NS crust may be close to reach thermal equilibrium with the core."," We detected only a decrease in flux and a decrease in temperature between the last two observations, indicating that the NS crust may be close to reach thermal equilibrium with the core."1176 Fitting the inferred temperatures with an exponential decay plus a constant offset yields an e-folding time of +888 days., Fitting the inferred temperatures with an exponential decay plus a constant offset yields an e-folding time of $\pm$ 88 days.1177" In addition to the thermal component, the first observation shows a component above ~2-3 keV that we fitted with a power-law component, and which represents ~77% of the total 0.3-10 keV flux."," In addition to the thermal component, the first observation shows a component above $\sim$ 2-3 keV that we fitted with a power-law component, and which represents $\sim$ of the total 0.3-10 keV flux."1178" The index of the power law is poorly constrained, with a value of 0.2443-0.7."," The index of the power law is poorly constrained, with a value of $\pm$ 0.7."1179" This value is much lower than the one found by ? in a previous analysis of this observation, 1.7+0.5, but the flux contribution to the total flux is similar."," This value is much lower than the one found by \citet{0748:degenaar10mnras} in a previous analysis of this observation, $\pm$ 0.5, but the flux contribution to the total flux is similar."1180 The index found in Obs 1 is consistent within the errors with values found for power-law components in other cooling NSs (e.g.?).., The index found in Obs 1 is consistent within the errors with values found for power-law components in other cooling NSs \citep[e.g.][]{cenx4:cackett10apj}.1181" Although the origin of this component is still unknown, it has been suggested that residual low-level accretion onto the magnetosphere or a shock from a pulsar wind could account for the flux levels observed (?).."," Although the origin of this component is still unknown, it has been suggested that residual low-level accretion onto the magnetosphere or a shock from a pulsar wind could account for the flux levels observed \citep{campana98aar}."1182 The fact that we do not detect dips in the light curves indicates that residual accretion onto the NS via an accretion disc is unlikely., The fact that we do not detect dips in the light curves indicates that residual accretion onto the NS via an accretion disc is unlikely.1183 We do not observe any significant contribution of the power-aw component after 6 November 2008., We do not observe any significant contribution of the power-law component after 6 November 2008.1184" We determined upper limits to the contribution of this component to the total flux of 1.1, 0.8 and in the oobservations on 17 March and 1 July 2009 and 17 June 2010, respectively."," We determined upper limits to the contribution of this component to the total flux of 1.1, 0.8 and in the observations on 17 March and 1 July 2009 and 17 June 2010, respectively."1185" In contrast, ? report a significant changing contribution between 5 and of the power-law component to the total flux based on Chandra observations of the source in 10 February and 5 June 2009 and in 15 April 2010 (note the proximity in time of the Chandra observations to the ones analysed in this work)."," In contrast, \citet{0748:degenaar10mnras} report a significant changing contribution between 5 and of the power-law component to the total flux based on $Chandra$ observations of the source in 10 February and 5 June 2009 and in 15 April 2010 (note the proximity in time of the $Chandra$ observations to the ones analysed in this work)."1186 'The sensitivity to detect and accurately model the power-law component is crucial to constrain the behaviour of the thermal component., The sensitivity to detect and accurately model the power-law component is crucial to constrain the behaviour of the thermal component.1187" In this sense, the higher effective area of ccompared to Chandra and the longer exposures make the observations presented in this work more suitable to determine the contribution of the power law to the total flux."," In this sense, the higher effective area of compared to $Chandra$ and the longer exposures make the observations presented in this work more suitable to determine the contribution of the power law to the total flux."1188" However, the different power law fluxes could be explained if the source varied between the Chandra and oobservations."," However, the different power law fluxes could be explained if the source varied between the $Chandra$ and observations."1189 We observe a significant decrease in the optical magnitude of the system., We observe a significant decrease in the optical magnitude of the system.1190 The most significant drop of intensity occurs after Obs 1 and is very small between Obs 2 and 4., The most significant drop of intensity occurs after Obs 1 and is very small between Obs 2 and 4.1191 This supports further the results from X-ray spectral fitting in the sense that any residual accretion present in Obs 1 has disappeared in subsequent observations., This supports further the results from X-ray spectral fitting in the sense that any residual accretion present in Obs 1 has disappeared in subsequent observations.1192 ? observed bbetween November 2008 and January 2009 using Andicam and the SMARTS 1.3 m telescope., \citet{0748:hynes09apj} observed between November 2008 and January 2009 using Andicam and the SMARTS 1.3 m telescope.1193" They found an average magnitude of R = 22.4 and J = 21.3 for the optical counterpart and a periodicity consistent with the orbital period of the system, indicating that at the time of the observations emission from the accretion disc and/or X-ray heated inner face of the companion star dominate the optical emission."," They found an average magnitude of R = 22.4 and J = 21.3 for the optical counterpart and a periodicity consistent with the orbital period of the system, indicating that at the time of the observations emission from the accretion disc and/or X-ray heated inner face of the companion star dominate the optical emission."1194" We do not observe any significant modulation in Obs 2 and 4, indicating that the emission from the accretion disc has most likely disappeared and the inner face of the companion star has subsequently cooled down."," We do not observe any significant modulation in Obs 2 and 4, indicating that the emission from the accretion disc has most likely disappeared and the inner face of the companion star has subsequently cooled down."1195" However, given the large errors of the measurements in Obs 2 and 4, the existence of a small modulation due to a remaining temperature gradient between the illuminated and dark face of the companion cannot be ruled out."," However, given the large errors of the measurements in Obs 2 and 4, the existence of a small modulation due to a remaining temperature gradient between the illuminated and dark face of the companion cannot be ruled out."1196 We attempted to constrain the value of the interstellar absorption with the high-resolution RGS spectra., We attempted to constrain the value of the interstellar absorption with the high-resolution RGS spectra.1197 This is important especially in the case that spectral variability in the power-law component is detected., This is important especially in the case that spectral variability in the power-law component is detected.1198 Coupled variations between the power-law index and the value of wwere found for Ααἱ X-1 and Cen X-4 (??) which could be explained if the power-law arises as shocked emission from the pulsar wind and the infalling material (?)..," Coupled variations between the power-law index and the value of were found for Aql X-1 and Cen X-4 \citep{aql:campana03apj,cenx4:campana04apj} which could be explained if the power-law arises as shocked emission from the pulsar wind and the infalling material \citep{campana98aar}. ."1199similar (ο that in voung stellar objects (e.g..Frank&Blackman2004).,"similar to that in young stellar objects \citep[e.g.,][]{fra04}."1200. There is growing evidence for binary interactions in AGB stars with high mass loss rates (e.g..Mauron&IInggins2006).. and (he binary. accretion disk scenario is consistent with our findings on several points. (," There is growing evidence for binary interactions in AGB stars with high mass loss rates \citep[e.g.,][]{mau06}, and the binary accretion disk scenario is consistent with our findings on several points. ("

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