ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 The most notable difference is related to dust which does uot affect the NIR images., The most notable difference is related to dust which does not affect the NIR images.3 Tidal tails and bridges seen as frequent in the NIR as in the optical, Tidal tails and bridges seem as frequent in the NIR as in the optical.4 Ilowever. we do observe a higher fraction of asvuunetric galaxies iu the NIR than iu the optical for the UCCs.," However, we do observe a higher fraction of asymmetric galaxies in the NIR than in the optical for the HCGs."5 Since we are not secius this effect iu the IKPGs. this clearly states that something is different in the IICCs.," Since we are not seeing this effect in the KPGs, this clearly states that something is different in the HCGs."6 Possibly the interactions in the IICCis are at a amore advanced stage and are affecting the oldest stellar population., Possibly the interactions in the HCGs are at a more advanced stage and are affecting the oldest stellar population.7 Or possibly. interactions aro now occurring d the absence of the usual evidences of star formation or nuclear activity.," Or possibly, interactions are now occurring in the absence of the usual evidences of star formation or nuclear activity."8 This last possibility is consistent with the dry merger hypothesis., This last possibility is consistent with the dry merger hypothesis.9 Tu normal early-type galaxies. color evadicuts make a galaxy core redder than the periphery (ic.. the eradient. is. negative:. Peleticr. οἳ αἱ.," In normal early-type galaxies, color gradients make a galaxy core redder than the periphery (i.e. the gradient is negative; Peletier et al."10 1990)., 1990).11 These color gradieuts can be explained. in part by the concentration of older stellar population towards the ceuter of the galaxies. and in other part. by an increase in stellar metallicity (IIlinklev&Tin32001).," These color gradients can be explained, in part by the concentration of older stellar population towards the center of the galaxies, and in other part, by an increase in stellar metallicity \citep{hinkley01}."12.. Galaxy formation models Sugeest that if au elliptical galaxy form rapidly lw monolithic collapse aud is undisturbed ly interaction. a negative color gradients will form aud stay uuchaneed for most of its life-time.," Galaxy formation models suggest that if an elliptical galaxy form rapidly by monolithic collapse and is undisturbed by interaction, a negative color gradients will form and stay unchanged for most of its life-time."13 However. some elliptical galaxies are known to present color gradients which are flat or positive. ie. bluer to the iuney part (Aichard1999:huctal.2001:Yanget 2006).," However, some elliptical galaxies are known to present color gradients which are flat or positive, i.e., bluer to the inner part \citep{michard99, im01, yang06}."14 For these ealaxies. nmodels sugeest that such features in color eradieuts cau be the result of mergers or past interactions with σαςτο] galaxies.," For these galaxies, models suggest that such features in color gradients can be the result of mergers or past interactions with gas-rich galaxies."15 As in Coziol αποπιΕναν (2007). we have search for NIR blue cores iu the early-type galaxies of our sunuple.," As in Coziol Plauchu-Frayn (2007), we have search for NIR blue cores in the early-type galaxies of our sample."16 The JA’ color gradieut is defined as δουν)., The $J - K'$ color gradient is defined as $\Delta (J-K')/log(r)$.17 According to this definition. ealaxies with blue cores have ACFA)flog(r)samples- 0.," According to this definition, galaxies with blue cores have $\Delta (J-K')/log(r)<0$ ."18 For 29 early-type galaxies in the three we were able to estimate this gradient., For 29 early-type galaxies in the three samples we were able to estimate this gradient.19" In these ealaxies we found colors consistent withblue cores or flat gradieuts in 10 out of 22 (15%)) Πο» aud bout of £()) τέως,", In these galaxies we found colors consistent with blue cores or flat gradients in 10 out of 22 ) HCGs and 4 out of 6 ) KPGs.20 The oulv carly-type, The only early-type21"also took into account the measurement by ?.272+10rad/m- From this set of five measurements we derived an RM of 255.01£0.83 rad/nv"".",also took into account the measurement by \citet[][$272\pm10$ $\mathrm{rad/m^2}$ From this set of five measurements we derived an RM of $255.01 \pm 0.83$ $\mathrm{rad/m^2}$.22" It should be clear that. with regard to the 350 MHz RM measurements. the fits give the same reduced y for both the positive and the negative correction to the initial ""derotation""."," It should be clear that, with regard to the 350 MHz RM measurements, the fits give the same reduced $\chi^2$ for both the positive and the negative correction to the initial ""derotation""."23 We removed those ambiguities by considering the Stokes U measurements near 850 MHz data on 2005 January 5., We removed those ambiguities by considering the Stokes U measurements near 850 MHz data on 2005 January 5.24 The fit to that data gave an RM of 253.14+12.43 rad/m which made all of the positive RM solutions to the 350 MHz data very unlikely (=3.0c level for January 4 and 7)., The fit to that data gave an RM of $253.14\pm12.43$ $\mathrm{rad/m^2}$ which made all of the positive RM solutions to the 350 MHz data very unlikely $\simeq 3.0 \sigma$ level for January 4 and 7).25 It is evident that the contribution of the ionosphere to the RM. RA; Is included in all fits.," It is evident that the contribution of the ionosphere to the RM, $\mathrm{RM_{ion}}$, is included in all fits."26" For the 2005 January 4. 5. 7 and 10 observations. RM;,, as reported by the AIPS task ""TECOR'. is the range 2.]+04 rad/m-."," For the 2005 January 4, 5, 7 and 10 observations, $\mathrm{RM_{ion}}$ as reported by the AIPS task 'TECOR', is the range $2.1\pm 0.4$ $\mathrm{rad/m^2}$."27 Consequently. the interstellar RM is given. by RMin=255.01-2.1252.91€0.92 rad/im-.," Consequently, the interstellar RM is given by $\mathrm{RM_{int}}=255.01-2.1=252.91 \pm 0.92$ $\mathrm{rad/m^2}$."28 We were able to measure the fractional linear. polarization on all of the four epochs mentioned in paragraph 3.2.1..., We were able to measure the fractional linear polarization on all of the four epochs mentioned in paragraph \ref{par:polgeneral}.29 At 850 MHz. we were not able to measure polarization on. 2005 January 10.," At 850 MHz, we were not able to measure polarization on 2005 January 10."30 For the other occasions. the measured polarized fluxes. P=4Q?-U2. fractions and their error bars are listed in table 4..," For the other occasions, the measured polarized fluxes, $P=\sqrt{Q^2+U^2}$, fractions and their error bars are listed in table \ref{tab:linpolfrac}."31 The latter two quantities are depicted in figure 3.., The latter two quantities are depicted in figure \ref{fig:linpolfracs}.32 The overall conclusion is that there is no compelling evidence for any significant depolarization at any frequency., The overall conclusion is that there is no compelling evidence for any significant depolarization at any frequency.33 Only the polarization fraction at 1300 MHz on January 4 is low compared to the 8.4 GHz measurements. but this fraction was determined from our worst fit. Le.. the fit with the highest reduced y. The polarization angles and their uncertainties are. also listed in table 4..," Only the polarization fraction at 1300 MHz on January 4 is low compared to the 8.4 GHz measurements, but this fraction was determined from our worst fit, i.e., the fit with the highest reduced $\chi^2$ The polarization angles and their uncertainties are also listed in table \ref{tab:linpolfrac}. ."34 The observations at 550 and 1300 MHz gave the most accurate position angles. with typical uncertainties of order 10°.," The observations at 850 and 1300 MHz gave the most accurate position angles, with typical uncertainties of order $10\degr$."35 They are depicted in figure 4.., They are depicted in figure \ref{fig:polangles}.36 Here. we see compelling evidence for significantly. different. polarization angles with respect to the 8.4 GHz observations from ?.. particularly on January 5 and 850 MHz and on January 10 at both 850 and 1300 MHz.," Here, we see compelling evidence for significantly different polarization angles with respect to the 8.4 GHz observations from \citet{Taylor2005}, particularly on January 5 and 850 MHz and on January 10 at both 850 and 1300 MHz."37 It is clear from figure 1. that SGRI806-20 is much dimmer at 350 MHz than what would be expected from the GMRT observations. at. 240 and 610 MHz (?).., It is clear from figure \ref{fig:fluxPband} that SGR1806-20 is much dimmer at 350 MHz than what would be expected from the GMRT observations at 240 and 610 MHz \citep{Cameron2005}.38 In principle the Luminous Blue Variable. 14 to the east of SGR 1806-20 (seetheSupplementaryInformationto?) should be easily distinguishable from the Soft Gamma Repeater in the GMRT images. even at 240 MHz.," In principle the Luminous Blue Variable, $14\arcsec$ to the east of SGR 1806-20 \citep[see the Supplementary Information to][]{Gaensler2005a} should be easily distinguishable from the Soft Gamma Repeater in the GMRT images, even at 240 MHz."39 The FWHM beamsize reported at that frequency is 12”«18 (?).., The FWHM beamsize reported at that frequency is $12\arcsec \times 18\arcsec$ \citep{Chandrab}.40 This makes it hard to understand the In principle the discrepancy cannot originate from. the inclusion or exclusion of extended emission., This makes it hard to understand the In principle the discrepancy cannot originate from the inclusion or exclusion of extended emission.41 The GMRT data were corrected for this (??)..," The GMRT data were corrected for this \citep{Chandraa,Chandrab}."42 We excluded short spacings («κ) from our 350 MHz WSRT observations., We excluded short spacings $< 1\mathrm{k}\lambda$ ) from our 350 MHz WSRT observations.43 This was actually a necessity since these were daytime observations and solar interference would otherwise compromise our calibration (seealso?.endofparagraph3.2).., This was actually a necessity since these were daytime observations and solar interference would otherwise compromise our calibration \citep[see also][end of paragraph 3.2]{Brentjens2008}.44 Also. it is possible that the LBV radio nebula is variable and that it was much brighter on 2005 April 30/May 1 than on some occasions in 2005 January.," Also, it is possible that the LBV radio nebula is variable and that it was much brighter on 2005 April 30/May 1 than on some occasions in 2005 January."45" We ran the AIPS task ""IMFIT' on the map from our 2005 April 30/May | observation and we found a peak flux density of 138+ImJy/beam and an integrated flux of 189+2mJy/beam at the location of the LBV.", We ran the AIPS task 'IMFIT' on the map from our 2005 April 30/May 1 observation and we found a peak flux density of $138\pm 1 \mathrm{mJy/beam}$ and an integrated flux of $189\pm 2 \mathrm{mJy/beam}$ at the location of the LBV.46 The NVSS (?) image of this field shows this source at the 15 mJy level., The NVSS \citep{Condon98} image of this field shows this source at the 15 mJy level.47 This would indicate that the LBV has a spectral index of about -1.8. which 1s almost the index for thermal radio radiation.," This would indicate that the LBV has a spectral index of about -1.8, which is almost the index for thermal radio radiation."48 It should be noted that. at the times of the latest observations in January 2005. when the radio nebula was relatively dim. there is no evidence for negative residuals in our maps that could be caused by the subtraction of the LBV.," It should be noted that, at the times of the latest observations in January 2005, when the radio nebula was relatively dim, there is no evidence for negative residuals in our maps that could be caused by the subtraction of the LBV."49 This indicates that. most likely. the LBV had the same brightness at the times of at least some of the 2005 January measurements as OI 2005 April 30/May 1.," This indicates that, most likely, the LBV had the same brightness at the times of at least some of the 2005 January measurements as on 2005 April 30/May 1."50 Variability at radio wavelengths of the radio nebulas from LBVs has been known for quite some time (see.e.g..?)..," Variability at radio wavelengths of the radio nebulas from LBVs has been known for quite some time \citep[see, e.g.,][]{Abbott81}."51 For the P Cygni nebula variability at timescales of days was established at em wavelengths (2)., For the P Cygni nebula variability at timescales of days was established at cm wavelengths \citep{Skinner96}.52 These authors report a 50% increase in flux in less than two days on one occasion during three months of observations on every other day., These authors report a $50\%$ increase in flux in less than two days on one occasion during three months of observations on every other day.53 It 1s unknown how these variations translate to lower frequencies., It is unknown how these variations translate to lower frequencies.54 We therefore cannot completely exclude that the LBV was brighter at the time of the 2005 April 30/May | observation than on some occasions in January 2005., We therefore cannot completely exclude that the LBV was brighter at the time of the 2005 April 30/May 1 observation than on some occasions in January 2005.55 Also. the spectral index derived above does not agree with any of the spectral indices of the four LBVs observed by ? at 3 and 6 em.," Also, the spectral index derived above does not agree with any of the spectral indices of the four LBVs observed by \citet{Duncan2002} at 3 and 6 cm."56 Two of those spectral indices are close to that of a spherically symmetric radially expanding stellar wind (+0.6.see??)..," Two of those spectral indices are close to that of a spherically symmetric radially expanding stellar wind \citep[+0.6, see][]{Panagia75,Wright75}."57 However. at these wavelengths. those systems may well be described as optically thin. which may not be the case at the frequencies we are The WSRT 850 MHz Stokes I measurements are not inconsistent. with the 840 MHz MOST data published earlier (?).. given the rather large noise levels in the data from both telescopes.," However, at these wavelengths, those systems may well be described as optically thin, which may not be the case at the frequencies we are The WSRT 850 MHz Stokes I measurements are not inconsistent with the 840 MHz MOST data published earlier \citep{Gaensler2005a}, given the rather large noise levels in the data from both telescopes."58 The last MOST observation was taken 15 days after the Giant Flare (GF)., The last MOST observation was taken 15 days after the Giant Flare (GF).59 Consequently. the 550 MHz WSRT observations after 2005 January 10 cannot be compared with other observations in this band.," Consequently, the 850 MHz WSRT observations after 2005 January 10 cannot be compared with other observations in this band."60 The last three of the January 2005 observations at 850 MHz were less contaminated by RFI than the first four. which resulted insmaller error bars on the fluxes.," The last three of the January 2005 observations at 850 MHz were less contaminated by RFI than the first four, which resulted insmaller error bars on the fluxes."61 There is evidence (>2c level) for a deviation from a power-law decay from about 15 days after the GF. analogous," There is evidence $>2\sigma$ level) for a deviation from a power-law decay from about 15 days after the GF, analogous"62Iu he rawe ds«d<20. the color-couuts are known to spli iuto two major peaks. cach sample a different stellar pxopulation (ασ 1986).,"In the range $18<V<20$, the color-counts are known to split into two major peaks, each sampling a different stellar population (Bahcall 1986)."63 The blue peak at (Bτο05 Is ¢ue to halo stars near the turuoff CM:~| 6). locate at few kpc from the Galactic plane.," The blue peak at $(B-V)\sim0.5$ is due to halo stars near the turnoff $M_V\sim+4$ ), located at few kpc from the Galactic plane."64 The red peak at (BVP)~1.3 is due to ant M-dwarf stars frou t10 cisk. locaed at less than 1 kpc frou the Sun.," The red peak at $(B-V)\sim1.3$ is due to faint M-dwarf stars from the disk, located at less than 1 kpc from the Sun."65 Note t10 sla] rolaive offset (0.1 mae) between he observed axd xedietecl location of the rec peak., Note the small relative offset $\sim0.1$ mag) between the observed and predicted location of the red peak.66 As ]xjuted out above his av )e due to. aud is οςyUSistCLL with. the departure oft1C οςlor term from the near οςTEEion adopted for Ojects with (5Wy)>12.," As pointed out above this may be due to, and is consistent with, the departure of the color term from the linear correction adopted for objects with $(B-V)>1.2$."67 Note tlat he agreement is uuch better in (WVOL) for which t1C οςntribution from color terms are expected to e negligible., Note that the agreement is much better in $(V-I)$ for which the contribution from color terms are expected to be negligible.68 Traditionally. the observed spitti negli the color peaks las Deel used to determine tιο local LOLIalization of halo stars in the solar ucighborlood.," Traditionally, the observed splitting in the color peaks has been used to determine the local normalization of halo stars in the solar neighborhood."69 Note in this context the difference in the amplitude of the counts iu the blue peal. in tje naeuitude range 202]. WwHeh ds seen iu both (BV) aud (VI).," Note in this context the difference in the amplitude of the counts in the blue peak in the magnitude range $20<V<21$, which is seen in both $(B-V)$ and $(V-I)$."70 Most photonetric surveys at faint magnitudes Reid and Majewski 1993) have relied on peucil-beam surveys covering a πια fraction of a degree., Most photometric surveys at faint magnitudes Reid and Majewski 1993) have relied on pencil-beam surveys covering a small fraction of a degree.71 Therefore. the uuuber of observed objects per biu has been very siunall. leading to large uncertainties in the derived model parameters.," Therefore, the number of observed objects per bin has been very small, leading to large uncertainties in the derived model parameters."72 The EIS suuple. covering 1.3 square degrees. represents a significant improvement auc may allow for a better determination of these paramcters.," The EIS sample, covering $\sim 1.3$ square degrees, represents a significant improvement and may allow for a better determination of these parameters."73 Finally. it is iuterestiug to poiut out the existence of a population of blue ojects. mn particuar. the sugeestionOO ofa peak at (DB.V)~(0.15 observed at faint maguitudes (20-V«21).," Finally, it is interesting to point out the existence of a population of blue objects, in particular, the suggestion of a peak at $(B-V)\sim0.15$ observed at faint magnitudes $<V<21$ )."74 This peak does not match the location zx the auplitude of the oeakk predicted by the white dwarf population asstumed iit 1ο model., This peak does not match the location and the amplitude of the peak predicted by the white dwarf population assumed in the model.75 Instead the observec blue objects could ««DISEN of a inix of white dwarts. blue horizoutal branch stars or perhays halo field bue stragelers.," Instead the observed blue objects could consist of a mix of white dwarfs, blue horizontal branch stars or perhaps halo field blue stragglers."76 Further invesieatiou on the nature of these objects ποστς wortlavhile., Further investigation on the nature of these objects seems worthwhile.77 Tιο results demonstrate that the stellar color cata[nn beige produced is by aid lurge consistent with model predictions and the observed differences may possibly point to deficiencies in the model which should be further investigated bv interested eroups., The results demonstrate that the stellar color catalog being produced is by and large consistent with model predictions and the observed differences may possibly point to deficiencies in the model which should be further investigated by interested groups.78 Altrough primarily cdaiven bv other goals. the above discussion shows that the EIS data is also useful for ealatic stidies.," Although primarily driven by other goals, the above discussion shows that the EIS data is also useful for galatic studies."79 Iu order to evaluate the quality aud the depth of the ealaxy ο galaxw counts in the differeut passbauds are shown in figure 15. and compared to those determined or patch A iud by other authors as indicated in the figure caption.," In order to evaluate the quality and the depth of the galaxy samples, galaxy counts in the different passbands are shown in figure \ref{ncounts_gal} and compared to those determined for patch A and by other authors as indicated in the figure caption."80 In these comparisons the Z maguities of Lidinan Peterson (1996) have heen shifted by |0.OL mae ancl hose measured by Postman (1996) 1 113 nae o bring them into the Johusou-Cousins systeui., In these comparisons the $I$ magnitudes of Lidman Peterson (1996) have been shifted by +0.04 mag and those measured by Postman (1996) by -0.43 mag to bring them into the Johnson-Cousins system.81 A small correction (-0.02 mae} has also been applied o the V counts of Postinaua£., A small correction (-0.02 mag) has also been applied to the $V$ counts of Postman.82.. No corrections were lade to he Arnouts (1997) data., No corrections were made to the Arnouts (1997) data.83 As can be seen here is a roinarkable agreement betwee- he EIS conts and those obtained by other authors., As can be seen there is a remarkable agreement between the EIS counts and those obtained by other authors.84 They are also cousiseut with he counts deteriuued from teh A. down to Vea and J~22.5.," They are also consistent with the counts determined from patch A, down to $V\sim24$ and $I\sim22.5$."85 As enmipliasize oe1 Paper I even or single exposures EIS reaches fainter magnitudes than previous data used for cluster searclies., As emphasized in Paper I even for single exposures EIS reaches fainter magnitudes than previous data used for cluster searches.86 The overall wuiformity of he EIS ealaxy catalogs cai be examined using the two-point angular correlation function. 60).," The overall uniformity of the EIS galaxy catalogs can be examined using the two-point angular correlation function, $w(\theta)$."87 Tudeed. w(@) is a very efficient tool for cleecting anv kind of artificial patterus (such as a erid WIh scale comparable to an EMMI frame) or possible eradicuts iu the deusity over the feld (which could resultfrom large-scale eracieuts of the photometric zero-point).," Indeed, $w(\theta)$ is a very efficient tool for detecting any kind of artificial patterns (such as a grid with scale comparable to an EMMI frame) or possible gradients in the density over the field (which could resultfrom large-scale gradients of the photometric zero-point)."88 Departures from wuiformity should affect the correlation function especially at faint magnuitudes., Departures from uniformity should affect the correlation function especially at faint magnitudes.89 Figure 16 shows w(0) obtained for cach of the three passbands B.V. aud. £. using the estimator proposed by Laudy," Figure \ref{fig:w} shows $w(\theta)$ obtained for each of the three passbands $B, V,$ and $I$ , using the estimator proposed by Landy"90This CME has the highest value of maximum acceleration of all (he CMEs studied here. which is over 1500mis.7 at a height of about ο...,"This CME has the highest value of maximum acceleration of all the CMEs studied here, which is over $1500\mpss$ at a height of about $2\Rsun$."91 Such a high value is observed for CATES associated with flares. which are termed as impulsive by Sheelevetal.(1999) and (2002).," Such a high value is observed for CMEs associated with flares, which are termed as impulsive by \citet{Sheeley.etal1999} and \citet{Moon.etal2002}."92. IIowever. the flare in this case was classified with X-ray class C8. and such hieh values of CME acceleration are earlier reported (o be associated wilh X-class flares.," However, the flare in this case was classified with X-ray class C8, and such high values of CME acceleration are earlier reported to be associated with X-class flares."93 Such a high value of acceleration has also been reported by Alexanderοἱal.(2002). for a CAE associated with an X1.2 class flare., Such a high value of acceleration has also been reported by \citet{Alexander.etal2002} for a CME associated with an X1.2 class flare.94 Also. assuming that this CME achieved its maximum value somewhere below 2Ro. we note that (this favours the CATE model proposed bv Chen&Ixrall(2003).. where (μον predict a bimodal acceleration profile.," Also, assuming that this CME achieved its maximum value somewhere below $2\Rsun$, we note that this favours the CME model proposed by \citet{Chen.Krall2003}, where they predict a bimodal acceleration profile."95 This CALE has been previously been analvsed by several researchers., This CME has been previously been analysed by several researchers.96 Among them. Temmeretal.(2010) have found acceleration of this CME to be 1300mis.7. while Linetal.(2010). have found it to be over 1000ms.7: both the results obtained [rom stereoscopic reconstruction of the CME.," Among them, \citet{Temmer.etal2010} have found acceleration of this CME to be $1300\mpss$, while \citet{Lin.etal2010} have found it to be over $1000\mpss$; both the results obtained from stereoscopic reconstruction of the CME."97 The large difference in acceleration values between our results. and those cited above. can be attributed to the different assumptions involved in the numerous reconstruction techniques (Mierlaetal.2010).," The large difference in acceleration values between our results, and those cited above, can be attributed to the different assumptions involved in the numerous reconstruction techniques \citep{Mierla.etal2010}."98. The CME on 2008 April 9 was associated with an active-region prominence. and was observed on the south-west solar limb. as shown in Figure 3..," The CME on 2008 April 9 was associated with an active-region prominence, and was observed on the south-west solar limb, as shown in Figure \ref{F:img09apr}."99 The CME first appeared in the AA and D FOVs at ULT and CUT respectively. while it could be just seen in COR2AÀ and D FOVs at UUT.," The CME first appeared in the A and B FOVs at UT and UT respectively, while it could be just seen in A and B FOVs at UT."100 The LE showed a bright knot close to its hiehest point. which was tracked during (he reconstruction.," The LE showed a bright knot close to its highest point, which was tracked during the reconstruction."101 The prominence material could be seen in images from 09:26 onwards in EUVIAA and ULT onwards in DD., The prominence material could be seen in images from 09:26 onwards in A and UT onwards in B.102the possible eruptiou survivor.,the possible eruption survivor.103" Their most likely candidate (""Object G). shown in Figure 1 of the presentpaper! has Vom 25.6. VRe L0. aud RO£z0.9 mae. roughly the colors of a mid-EK-tvpe supergiaut: if cereddenecd by reasonable amounts of interstellar and cieuiustellur reddening. the colors are cousistent with those of O-tvpe stars."," Their most likely candidate (“Object 6”), shown in Figure 1 of the present, has $V \approx 25.6$ , $V-R \approx 1.0$ , and $R-I \approx 0.9$ mag, roughly the colors of a mid-K-type supergiant; if dereddened by reasonable amounts of interstellar and circumstellar reddening, the colors are consistent with those of O-type stars."104 F95 point out that unaberrated miages. particularly in bluer bands. are required to better isolate the possible survivor.," F95 point out that unaberrated images, particularly in bluer bands, are required to better isolate the possible survivor."105 Based on its fading radio enmuüssiou recently detected iu new observations using the Verv Large Array (ΝΤΑ). Stockdale et al. (," Based on its fading radio emission recently detected in new observations using the Very Large Array (VLA), Stockdale et al. ("1062001) argue that the current radio properties of SN 1961V are consistent with those of some “peculiar Type II radio SNe.,2001) argue that the current radio properties of SN 1961V are consistent with those of some “peculiar” Type II radio SNe.107 Iu. particular. thev compare SN 1961V. with SN 1986J in NGC sol (Weiler. Panagia. Sraimels 1990). which was originally classified as Type V (Rupen et al.," In particular, they compare SN 1961V with SN 1986J in NGC 891 (Weiler, Panagia, Sramek 1990), which was originally classified as Type V (Rupen et al."108 1987). but is geucrally considered to be a prototvpical Type πι SN (see Schlegel 1990 and Filippeuko 1997 for discussion of this SN subtype).," 1987), but is generally considered to be a prototypical Type IIn SN (see Schlegel 1990 and Filippenko 1997 for discussion of this SN subtype)."109 Stockdale et al., Stockdale et al.110 coutend that it is less likely. based on its radio properties. that SN 1961V was similar to 5 Car or other LBVs. (," contend that it is less likely, based on its radio properties, that SN 1961V was similar to $\eta$ Car or other LBVs. ("111However. there is a lack of radio observatious of bright LDVs with ages similar to that of SN 1961V. preventing a secure determination of the true nature of SN 1961V.) Iu other words. they argue there should be no survivor. only a very old SN or voung SN remnant clucreine forty vears after explosion.,"However, there is a lack of radio observations of bright LBVs with ages similar to that of SN 1961V, preventing a secure determination of the true nature of SN 1961V.) In other words, they argue there should be no survivor, only a very old SN or young SN remnant emerging forty years after explosion."112 Iu this paper we further investigate whether SN LOGLA\ survived a LBV superoutburst by exploiting archivalHIST WEPC2 nuages that contain the SN site., In this paper we further investigate whether SN 1961V survived a LBV superoutburst by exploiting archival WFPC2 images that contain the SN site.113" We recover ""Object 6"" of F95. aud we show that an fait and very red object is consistent with the radio position of SN 1961V aud may represent the possible surviving star."," We recover “Object 6” of F95, and we show that an faint and very red object is consistent with the radio position of SN 1961V and may represent the possible surviving star."114 We consider it less likely to be a facing. very old SN. but additional observations are necessary to further clarity this ISSTIC.," We consider it less likely to be a fading, very old SN, but additional observations are necessary to further clarify this issue."115 TwoHST programs. CO-5116 and CO-9012. have iuaged NGC 1058. the host ealaxy of ον 1961V: the datasets are publicly available in the archive.," Two programs, GO-5446 and GO-9042, have imaged NGC 1058, the host galaxy of SN 1961V; the datasets are publicly available in the archive."116 The former program obtained ou 1991 September 8 a pair of NÜ-« exposures with the F606W filter. while the latter prograun obtained on 2001 July 23 two pairs of 230-s exposures. one pair with the F150W filter and the other with the ESIIW filter.," The former program obtained on 1994 September 8 a pair of 80-s exposures with the F606W filter, while the latter program obtained on 2001 July 3 two pairs of 230-s exposures, one pair with the F450W filter and the other with the F814W filter."117 These bandpasses roughly correspond to V. 2. aud 7.," These bandpasses roughly correspond to $V$, $B$, and $I$."118 Using the radio position of SN 1961V from Stockdale et al. (, Using the radio position of SN 1961V from Stockdale et al. (1192001: it is coincident with the position measured by Cowan ct al.,2001; it is coincident with the position measured by Cowan et al.120 1988. to within the errors). aud correcting for geometric distortions on the WES chip. we isolate the site of the SN for both sets of observations.," 1988, to within the errors), and correcting for geometric distortions on the WF3 chip, we isolate the site of the SN for both sets of observations."121 Figures 2 and 3 show the SN site ou the FSLIW aud F150WNV images. respectively.," Figures 2 and 3 show the SN site on the F814W and F450W images, respectively."122" An error circle is plotted at the radio position. based ou the astrometric itormation in the headers of the FLISOW and FaliW nuages. with a (conservative) radius of ~175 (the observatious by GO-9012 ciuploved the fine-lock ποσο, which results in about this level of astrometric macertaity: see the discussions ofHST astrometry in F95 aud Van Dyk et al."," An error circle is plotted at the radio position, based on the astrometric information in the headers of the F450W and F814W images, with a (conservative) radius of $\sim 1{\farcs}5$ (the observations by GO-9042 employed the fine-lock mode, which results in about this level of astrometric uncertainty; see the discussions of astrometry in F95 and Van Dyk et al."123 19995)., 1999b).124" Figure | shows the site in the F6OOGW nage: this exposure was obtained in ""evro mode and likely has less reliable astrometry than the ΕΣΑΝ aud FLOW images. so we inpose the error circle of Figures 2 and 3 onto Figure 1."," Figure 4 shows the site in the F606W image; this exposure was obtained in “gyro” mode and likely has less reliable astrometry than the F814W and F450W images, so we impose the error circle of Figures 2 and 3 onto Figure 4."125 The source numbering is from Figure 1. as in F95.," The source numbering is from Figure 1, as in F95."126 Object Sis situated off the ΕΟΝ and FsliW images. but is detected ou the F6OGW tage.," Object 8 is situated off the F450W and F814W images, but is detected on the F606W image."127 To measure the brightuesses of the various sources hrough each of the three bands. we employed version 1.1 of the package IISTphot (Dolphin 2000a.b).," To measure the brightnesses of the various sources through each of the three bands, we employed version 1.1 of the package HSTphot (Dolphin 2000a,b)."128" We followed he ""recipe? iu the HSTphot maul. initially adopting a la detection threshold aud using the following asks in sequeutial order:eeskcrmeshk.coadd. gotshg. hotpirels. and hstphot."," We followed the “recipe” in the HSTphot manual, initially adopting a $4\sigma$ detection threshold and using the following tasks in sequential order:, and ."129 Ax Dolphin (2000b) has shown. IISTphot produces results that are quite consistent with roxe obtained from both the DAOPIIOT and DoPIIOT oickases. while accounting for ΝΕΟΣ poiut-spread Muction variations and charge-transfer effects across the ‘hips. zeropoiuts. aperture corrections. etc..," As Dolphin (2000b) has shown, HSTphot produces results that are quite consistent with those obtained from both the DAOPHOT and DoPHOT packages, while accounting for WFPC2 point-spread function variations and charge-transfer effects across the chips, zeropoints, aperture corrections, etc.,"130 automatically within one package. (, automatically within one package. (131We have also conducted Linited tests of IISTphot DAOPIIOT aud find very good aerecment neneithe results.),We have also conducted limited tests of HSTphot DAOPHOT and find very good agreement in the results.)132 Table 1 eives the results of our photonietry i the flight «παν baudpasses., Table 1 gives the results of our photometry in the flight system bandpasses.133 Not all sources were etected in all three bands. given the relatively low signal-o-nolse ratio (S/N) of these images.," Not all sources were detected in all three bands, given the relatively low signal-to-noise ratio (S/N) of these images."134 Detection limits (20) ave Mpiso©25.3. meu©25.L sud esppym25.3 nag.," Detection limits $3\sigma$ ) are $m_{\rm F450W} \approx 25.3$, $m_{\rm F606W} \approx 25.4$, and $m_{\rm F814W} \approx 25.3$ mag."135 Iu order to further analyze the archival inages. we lave ransformied our flight system magnitudes into Joliuson-Cousins BV imaguitudes. in an analogous mamnucr to hat of F95. i.0.. via svuthetic photometry of normal stars of a wide span of spectral types and luminosity classes. obtained by applying the STSDAS package SYNPIIOT to he Bruzual Spectral Svuthetic Atlas.," In order to further analyze the archival images, we have transformed our flight system magnitudes into Johnson-Cousins $BVI$ magnitudes, in an analogous manner to that of F95, i.e., via synthetic photometry of normal stars of a wide span of spectral types and luminosity classes, obtained by applying the STSDAS package SYNPHOT to the Bruzual Spectral Synthetic Atlas."136 In Table 1 we list he BVT magnitudes. wherever possible. for the objects in the SN 1961V. euviromneut.," In Table 1 we list the $BVI$ magnitudes, wherever possible, for the objects in the SN 1961V environment."137 Iu two cases; woe used the detection limit at P606W to set a lower limit ou the 7 uaenitude: for Object 8. we could only approximate the ranstormation.," In two cases, we used the detection limit at F606W to set a lower limit on the $I$ magnitude; for Object 8, we could only approximate the transformation."138 Despite the passband differcuces between the two studies (P606 and FalWW here E555W and F7S5SLP in F95). the V. and J maenitudes of Objects 110 aeree airly well overall.," Despite the passband differences between the two studies (F606W and F814W here F555W and F785LP in F95), the $V$ and $I$ magnitudes of Objects 1–10 agree fairly well overall."139 Our V-band magnitudes are brighter wo O08 mae. but with a rather laree dispersion. 0.31 nae.," Our $V$ -band magnitudes are brighter by $-0.08$ mag, but with a rather large dispersion, 0.34 mag."140 Our L-band magnitudes are fainter by 0.01 mae. with a dispersion of 0.20 mae.," Our $I$ -band magnitudes are fainter by 0.04 mag, with a dispersion of 0.20 mag."141 The average errors iu our V and J magnitudes are 0.18 and 0.17 mae. respectively: or the measurements iu E95. they are 0.11 and 0.20 mag. respectively.," The average errors in our $V$ and $I$ magnitudes are 0.18 and 0.17 mag, respectively; for the measurements in F95, they are 0.14 and 0.20 mag, respectively."142 Thus. the Z-baud magnitudes agree to witlin he errors. but the dispersion iu our V magnitudes is about a factor of two larecr than the uncertainties in cither study.," Thus, the $I$ -band magnitudes agree to within the errors, but the dispersion in our $V$ magnitudes is about a factor of two larger than the uncertainties in either study."143 We cannot account for this discrepancy eutirelv hrough differcuces in bandpass or S/N: it may arise from underestimates of the unucertaintv im the aberrated WEPC photometry., We cannot account for this discrepancy entirely through differences in bandpass or S/N; it may arise from underestimates of the uncertainty in the aberrated WFPC photometry.144 For example. the largest disagreecnieuts are for Object 5 Gvhlich is clearly exteuded). Object 7 (πο is in a crowded environment withahigh backeround). aud Object LO Gvhich also may well be extended).," For example, the largest disagreements are for Object 5 (which is clearly extended), Object 7 (which is in a crowded environment withahigh background), and Object 10 (which also may well be extended)."145 Iu Figures 5 and 6 we show the color-magnitude diagrams for the objects in the SN LOGLV enviroment., In Figures 5 and 6 we show the color-magnitude diagrams for the objects in the SN 1961V environment.146 As, As147Using the 47 fit statistic. which accounts only for errors in the Pigs values. we find cf=0.67x0.07 and D=T4d 0.08. with o53 for T degrees of freedom. (the quoted uüncertainties on οἱ and. 2 are GS per cent confidence intervals).,"Using the $\chi^2$ fit statistic, which accounts only for errors in the $P_{\rm Bondi}$ values, we find $A=0.67\pm0.07$ and $B=0.74\pm0.08$ , with $\chi^2=5.3$ for 7 degrees of freedom (the quoted uncertainties on $A$ and $B$ are 68 per cent confidence intervals)."148 Using the BCES(Y LY) estimator of Akritas Bershacly (1996). which accounts for errors in both axes and he presence of intrinsic scatter. we obtain cl=V64£0.08S and B=pose=O0.78S£0.15.," Using the $Y|X$ ) estimator of Akritas Bershady (1996), which accounts for errors in both axes and the presence of possible intrinsic scatter, we obtain $A=0.64 \pm 0.08$ and $B=0.78 \pm0.15$."149"s The mean wean about he best fitting X moclelis στου£.)""VL."," The mean deviation about the best fitting BCES model is $\sigma({\rm log}\,P_{\rm Bondi})=0.11$."150 The results ο nel and D quoted above not account or the ellects. «X intrinsic (svstematic) scatter in the Maw loge relation., The results on $A$ and $B$ quoted above do not account for the effects of intrinsic (systematic) scatter in the $M_{\rm BH}-$ $\sigma$ relation.151 We have used further Monte Carlo simulations to examine these elfeets. introducing an intrinsic dispersion. of 146 dex in Puoi.," We have used further Monte Carlo simulations to examine these effects, introducing an intrinsic dispersion of 0.46 dex in $P_{\rm Bondi}$."152 Fits to the simulated data sets using the ICESQ'[X) estimator give A=0.65d0.16 and 5above)TrXx0220. in good agreement with the results presented but with (slightly) larger error bars on the fit parameters.," Fits to the simulated data sets using the $Y|X$ ) estimator give $A=0.65\pm0.16$ and $B=0.77\pm0.20$, in good agreement with the results presented above but with (slightly) larger error bars on the fit parameters."153 The results shown in Fig indicate the presence of a strong correlation between Za;4 and. {οι , The results shown in Fig \ref{fig:pbondi} indicate the presence of a strong correlation between $P_{\rm Bondi}$ and $P_{\rm jet}$.154A power-law mocel provides a good description of the data., A power-law model provides a good description of the data.155 Llowever. in eauging the origin and significance of this correlation. we must also consider clleets that could. arise from the plotted quantities having factors in common.," However, in gauging the origin and significance of this correlation, we must also consider effects that could arise from the plotted quantities having factors in common."156 Both £a; and. P4 depend on the distances to the objects., Both $P_{\rm Bondi}$ and $P_{\rm jet}$ depend on the distances to the objects.157 However. PosiXdiUL 7. whereas Dy.(Oxdio.LA c," However, $P_{\rm Bondi}\propto d_{\rm L}^{-0.5}$ , whereas $P_{\rm jet}\propto d_{\rm L}^{1.5}$."158pThus. the distance. dependences cannot lead to the observed. positive correlation.," Thus, the distance dependences cannot lead to the observed positive correlation."159 Both axes in Fig 4. also depend on the temperature. Z7 of the X- emitting gas., Both axes in Fig \ref{fig:pbondi} also depend on the temperature $T$ of the X-ray emitting gas.160" llowever. the temperature shows little variation [rom oect to object ancl varies only mildly with raclius between the! accretion and bubble μασ,"," However, the temperature shows little variation from object to object and varies only mildly with radius between the accretion and bubble radii."161 Moreover. the dependences on temperature are approximately suucxT57 and PixTt. which cannot lead to the observed positive correlation.," Moreover, the dependences on temperature are approximately $P_{\rm Bondi}\propto T^{-1.5}$ and $P_{\rm162jet}\propto T^{1.5}$, which cannot lead to the observed positive correlation."163" Finally. both axes involve the gas ensitv. measured at the accretion racius for 5,44; and at 10 bubble centres for 75.4."," Finally, both axes involve the gas density, measured at the accretion radius for $P_{\rm Bondi}$ and at the bubble centres for $P_{\rm jet}$."164 Phese radii are very dillerent ancl 1e density. profiles vary. significantly from object to object. meaning that the densities are essentially uncorrelated. (," These radii are very different and the density profiles vary significantly from object to object, meaning that the densities are essentially uncorrelated. ("165La detail. a mild anti-correlation is observed).,"In detail, a mild anti-correlation is observed)."166 We conclude that the observed. correlation. between Poon and Pa1 indicates a tightIn physical connection between je wo quantities., We conclude that the observed correlation between $P_{\rm Bondi}$ and $P_{\rm jet}$ indicates a tight physical connection between the two quantities.167" We have shown that [ου supermassive black holes at the centres of large. N-rav. luminous elliptical galaxies. a remarkable. tight correlation exists between the ""Bondi accretion rates inferred from. the Chandra X-ray cata ancl observed. galaxy velocity dispersions. and the power emerging [rom these systems in relativistic jets."," We have shown that for supermassive black holes at the centres of large, X-ray luminous elliptical galaxies, a remarkable, tight correlation exists between the `Bondi' accretion rates inferred from the Chandra X-ray data and observed galaxy velocity dispersions, and the power emerging from these systems in relativistic jets."168 Our result jt important implications for the nature ofthe accretion oocess and for issues relating to feedback. the growth. of lack holes and galaxy. formation.," Our result has important implications for the nature ofthe accretion process and for issues relating to feedback, the growth of black holes and galaxy formation."169 The relationship between the Bondi accretion power andjot power can be described bv a MN law model ofthe orm logμα=0.65(2E0.16)|0.77(0.20)log £34. where Επ and Poor is the power associated with inflating the cavities ancl providing the internal cnerev of he plasma that fills them.," The relationship between the Bondi accretion power and jet power can be described by a power law model of the form $P_{\rm Bondi} =1700.65(\pm0.16) + 0.77(\pm0.20)$ $\,P_{\rm jet}$, where $P_{\rm171Bondi} = 0.1 {\dot M_{\rm Bondi}} c^2$ and $P_{\rm jet}$ is the power associated with inflating the cavities and providing the internal energy of the plasma that fills them."172 Asiegnificant fraction (2.2- cent. for Po=107CODES J of the energy associated with he rest mass of material entering the Bondi accretion radius emerges from the systems in relativistic jets.," A significant fraction $2.2^{+1.0}_{-0.7}$per cent, for $P_{\rm jet}=10^{43}$ ) of the energy associated with the rest mass of material entering the Bondi accretion radius emerges from the systems in relativistic jets."173" Fhere is a slight indication that this fraction increases as 2, rises. [rom àl0 ↓⋅∆∫≻∩⊽⊓↓≻⋖⋅↓⋅⊓⋅⊔∣⋜∐∫↗↰∣⇁↕EU ⋅− ⋡∣∪↶∫≻⋅⋀⊥≓↓≻⋖⊾↓⋅⊓⊾⊔⋜⊔ ∫↗↰∣⊽↕∶1034 ("," There is a slight indication that this fraction increases as $P_{\rm jet}$ rises, from $1.3^{+1.0}_{-0.6}$ per cent at $P_{\rm174jet}=10^{42}$ , to $3.7^{+3.3}_{-1.7}$ per cent at $P_{\rm175jet}=10^{44}$ . ("176Llere. the quoted uncertainties include all sources of statistical and svstematic error discussed. in,"Here, the quoted uncertainties include all sources of statistical and systematic error discussed in"177Vhe X-ray. binary N2127|119. (AC211) in the core of the globular cluster MIS is optically one of the brightest and most well-stuclied of the low-mass X-ray binaries (LAINBs). and vet it remains a highly enigmatic svstem.,"The X-ray binary X2127+119 (AC211) in the core of the globular cluster M15 is optically one of the brightest and most well-studied of the low-mass X-ray binaries (LMXBs), and yet it remains a highly enigmatic system."178 Is optical and X-ray light-curves. along with its very low Lx/La4. make it a classic accretion cise corona (ADC) source (Fabian. CGuilbert Callanan 1987: Navlor et al.," Its optical and X-ray light-curves, along with its very low $L_{\rm X}$ $L_{\rm opt}$, make it a classic accretion disc corona (ADC) source (Fabian, Guilbert Callanan 1987; Naylor et al."179 1988). in which the system is seen almost edge-on and the compact object and hot. luminous inner disc are obscured by the accretion disc rim.," 1988), in which the system is seen almost edge-on and the compact object and hot, luminous inner disc are obscured by the accretion disc rim."180 Ehe X-ray [lux we observe comes entirely from photon-scattering by a Large corona above the disc. and is only a small fraction of the source's intrinsic X-ray Εις.," The X-ray flux we observe comes entirely from photon-scattering by a large corona above the disc, and is only a small fraction of the source's intrinsic X-ray flux."181 The N-rav source X2127|119. has. recently been discovered by to be two separate sources (White Angelini 2001)., The X-ray source X2127+119 has recently been discovered by to be two separate sources (White Angelini 2001).182 The second X-ray binary is 2.7 aresec from ο11. is 4 magnitudes fainter in the U-band.. but its count-rate is 2.5 times hieher.," The second X-ray binary is 2.7 arcsec from AC211, is 4 magnitudes fainter in the U-band, but its count-rate is 2.5 times higher."183 The discovery of a second. LAINB in the core of AILS explains one of most puzzling aspects of C211: luminous X-ray bursts showing expansion of the neutron star photosphere have xen observed in N21271119 (Dotani et al., The discovery of a second LMXB in the core of M15 explains one of most puzzling aspects of AC211: luminous X-ray bursts showing expansion of the neutron star photosphere have been observed in X2127+119 (Dotani et al.184 19080. van raraclijs et al.," 1990, van Paradijs et al."185 1990. Smale 2001). which has been dilficult to reconcile with AC?11's optical and X-ray light curves which showed it to be an ADC to observe N-rav. bursts with shotospheric expansion the neutron star surface has to be visible. but in ADC sources the neutron star is hidden from view by the aceretion disc.," 1990, Smale 2001), which has been difficult to reconcile with AC211's optical and X-ray light curves which showed it to be an ADC – to observe X-ray bursts with photospheric expansion the neutron star surface has to be visible, but in ADC sources the neutron star is hidden from view by the accretion disc."186 Fhis problem goes away if it is he second LAINB. and not AC211I. which is the burster.," This problem goes away if it is the second LMXB, and not AC211, which is the burster."187 The discovery may solve one mystery but. it produces another: the fact that ΑςΗν X-ray. luminosity is even ainter than previously believed. makes its unusually high optical luminosity even more puzzling. and suggests that the central X-ray source hidden from view must be exceptionally uminous. possibly indicating a very high mass-transler rate rom the companion star.," The discovery may solve one mystery but it produces another: the fact that AC211's X-ray luminosity is even fainter than previously believed makes its unusually high optical luminosity even more puzzling, and suggests that the central X-ray source hidden from view must be exceptionally luminous, possibly indicating a very high mass-transfer rate from the companion star."188 Determining whether AC211 has an extremely high nmiass-transfoer rate and is in an unusual evolutionary state is important: it is relevant to the understanding of LAINB evolution and to the understanding of stellar interactions within. and the evolution of. globular clusters.," Determining whether AC211 has an extremely high mass-transfer rate and is in an unusual evolutionary state is important: it is relevant to the understanding of LMXB evolution and to the understanding of stellar interactions within, and the evolution of, globular clusters."189 Theoretical determinations of the numbers of neutron stars in globular clusters ancl the ellicieney with which they interact with stars in the cluster cores to form binaries. combined with the large number of millisecond radio pulsars (end-products of LAINB evolution) observed in the clusters. imply that we see far fewer LAINBs in elobular clusters than we should.," Theoretical determinations of the numbers of neutron stars in globular clusters and the efficiency with which they interact with stars in the cluster cores to form binaries, combined with the large number of millisecond radio pulsars (end-products of LMXB evolution) observed in the clusters, imply that we see far fewer LMXBs in globular clusters than we should."190 This, This191backeround.,background.192 Ignoring the weak dependence of // and £ on e. appropriate for a hare binary. we find as before 14," Ignoring the weak dependence of $H$ and $L$ on $a$, appropriate for a hard binary, we find as before (a) = The exponent in this expression is of order $\mf/\m12\ll1$."193" where j4, has been set to unity. corresponding to an initial orientation parallel to the 0— axis."," Hence we can write 1 + ) where $\overline{\mu}_0$ has been set to unity, corresponding to an initial orientation parallel to the $\theta=0$ axis."194" For the case of binary supermassive black holes. if we define ej as the separation when the binary first forms a bound pair. we expect gravitational radiation coalescence to occur when af,210 (Merritt2000)."," For the case of binary supermassive black holes, if we define $a_0$ as the separation when the binary first forms a bound pair, we expect gravitational radiation coalescence to occur when $a/a_0\approx 10^{-2}$ \citep{mer00}."195. Since L/2//[cz2 for a hard binary. we find [or the expectation value of µ αἱ coalescence: a," Since $L/2H\approx 2$ for a hard binary, we find for the expectation value of $\mu$ at coalescence: 1 -."196pprorl —Mis Wriling 00=\/2(1—77). the rms change in the angle defined by the binarys spin axis. (his becomes," Writing $\delta\theta \equiv \sqrt{2(1-\overline{\mu})}$, the rms change in the angle defined by the binary's spin axis, this becomes"197eiven bx the neutron decoupling line in the figure.,given by the neutron decoupling line in the figure.198 If the flow is neutron rich. neutron decoupling will result in high energy (several hundred MeV) neutron-nucleus collisions whieh will destroy and svuthesize nuclei.," If the flow is neutron rich, neutron decoupling will result in high energy (several hundred MeV) neutron-nucleus collisions which will destroy and synthesize nuclei."199 These non-thermal reactions are nol described by the Kawano reaction network., These non-thermal reactions are not described by the Kawano reaction network.200 The net result of these collisions depends on the nuclear composition at the time of decoupling., The net result of these collisions depends on the nuclear composition at the time of decoupling.201 We will consider (wo instructive limiting cases: (1) deuteron rich and alpha poor. and (ii) deuteron poor and alpha rich.," We will consider two instructive limiting cases: (i) deuteron rich and alpha poor, and (ii) deuteron poor and alpha rich."202 Case (1) occurs for short dynamic timescales and [ον entropies well above the neutron decoupling line. while case (ii) occurs lor longer dynamic (ünmescales and entropies near the neutron decoupling line.," Case (i) occurs for short dynamic timescales and for entropies well above the neutron decoupling line, while case (ii) occurs for longer dynamic timescales and entropies near the neutron decoupling line."203 The final nucleosvnthesis is also sensitive to the (unknown) details of neutron. decoupling., The final nucleosynthesis is also sensitive to the (unknown) details of neutron decoupling.204 In. particular. we will see Chat the final deuterium abundance is sensitive to the exact number of high energv interactions suffered bv the average nucleus.," In particular, we will see that the final deuterium abundance is sensitive to the exact number of high energy interactions suffered by the average nucleus."205 This number is certainly between one and a few. but the exact number can be obtained only wilh a detailed transport calculation.," This number is certainly between one and a few, but the exact number can be obtained only with a detailed transport calculation."206 Anv deuterons present when neutron decoupling occurs will be destroved., Any deuterons present when neutron decoupling occurs will be destroyed.207 Likewise. anv a particles present will be broken apart.," Likewise, any $\alpha$ particles present will be broken apart."208 Because the branching ratio for a+(~GeV. is high. about 50% (Dimopoulosetal.1988).. the production of ?II through the spallation of às max (over)compensate the loss of deuterium in direct collisions.," Because the branching ratio for $\alpha+({\rm \sim209GeV\,\,\,nucleon}) \rightarrow {^2}{\rm H}+X$ is high, about $50\%$ \citep{dim}, the production of $\hh$ through the spallation of $\alpha$ 's may (over)compensate the loss of deuterium in direct collisions."210 Certainly for case (i). the a poor case. neutron decouplineg will result in à net decrease in deuterium.," Certainly for case (i), the $\alpha$ poor case, neutron decoupling will result in a net decrease in deuterium."211" For case (1), where a particles dominate. neutron decoupling will drive Yj to approximately half of the initial (i.e. at the beginning of neutron decoupling) value of Yay. provided that the average nucleus only undergoes one high energv collision."," For case (ii), where $\alpha$ particles dominate, neutron decoupling will drive $Y_{\rm D}$ to approximately half of the initial (i.e. at the beginning of neutron decoupling) value of $Y_{{^4}\rm He}$, provided that the average nucleus only undergoes one high energy collision."212 If the average nucleus undergoes a few high energy collisions. the nucleosvuthesis is somewhere inbe(ween the one collision case and the fixed point case discussed by Dimopoulosetal.(1988).," If the average nucleus undergoes a few high energy collisions, the nucleosynthesis is somewhere inbetween the one collision case and the fixed point case discussed by \cite{dim}."213 For example. if initially the material is composed entirely of a particles and [ree neutrons. and each a particle suffers one high energv collision with a neutron. then ~1054 of the material will be converted to ?IL.," For example, if initially the material is composed entirely of $\alpha$ particles and free neutrons, and each $\alpha$ particle suffers one high energy collision with a neutron, then $\sim 10\%$ of the material will be converted to $\hh$."214 Regions in Figure 2. above the decoupling line where the Ireeze-out mass Iraction of α particles is large could approximate this case., Regions in Figure \ref{fourlines} above the decoupling line where the freeze-out mass fraction of $\alpha$ particles is large could approximate this case.215 If additionally (his initial spallation takes place right at (he neutron decoupling point. (hen none of the deuterium produced from the fragmenüng of a will be destroyed.," If additionally this initial spallation takes place right at the neutron decoupling point, then none of the deuterium produced from the fragmenting of $\alpha$ will be destroyed."216 Only if these conditions are met can the final ?II vield be as high as ~LOW., Only if these conditions are met can the final $\hh$ yield be as high as $\sim10\%$.217 Note that here there is no source of significant post-nucleosvntliesis photo-dissociation ol deuterium., Note that here there is no source of significant post-nucleosynthesis photo-dissociation of deuterium.218 This is because on average there will only be LOOMeV per barvon in the form of electromagnetic radiation coming from the decay of pious produced in inelastic nucleon-nucleon collisions., This is because on average there will only be $\sim 100 {\rm MeV}$ per baryon in the form of electromagnetic radiation coming from the decay of pions produced in inelastic nucleon-nucleon collisions.219 This is three orders of magnitude smaller than the L00GeV per barvon coming trom the decay of massive particles discussed by Dimopoulos et al.. In, This is three orders of magnitude smaller than the $\sim 100{\rm GeV}$ per baryon coming from the decay of massive particles discussed by Dimopoulos et al.. In2201996).,.221. The SBF technique is based ou (he spectral analvsis of an image., The SBF technique is based on the spectral analysis of an image.222 Because the power spectrum of the image is modilied by the drizzling process. drizzled images were nol considered in this study. “," Because the power spectrum of the image is modified by the drizzling process, drizzled images were not considered in this study. “"223Weighted and cosmic-ray cleaned” images have been used instead.,Weighted and cosmic-ray cleaned” images have been used instead.224 These are fIat-fielded. cosmic-rav-rejected and skv-subtracted stacked images at each dither position.," These are flat-fielded, cosmic-ray-rejected and sky-subtracted stacked images at each dither position."225 Only dark exposures have been selected., Only dark exposures have been selected.226 Each of the the three wide-field ancl the planetary camera WEDPC? chips were analvzed separately., Each of the the three wide-field and the planetary camera WFPC2 chips were analyzed separately.227 The labels of all images considered. the corresponding filter and the total exposure times are listed in table I..," The labels of all images considered, the corresponding filter and the total exposure times are listed in table \ref{t-data}."228 The SBF concept was introduced by Tonry&Schneider(1988).. who noted that. in the surface photometry of a galaxy lar enough away (o remain unresolved. a fluctuation is observed because of the Poisson statistics of (he spatial distribution of stars. globular clusters. background. galaxies. etc.," The SBF concept was introduced by \citet{TS88}, who noted that, in the surface photometry of a galaxy far enough away to remain unresolved, a pixel-to-pixel fluctuation is observed because of the Poisson statistics of the spatial distribution of stars, globular clusters, background galaxies, etc."229 This technique was introduced. with the aim of measuring distances., This technique was introduced with the aim of measuring distances.230 Comparing SBFs produced by the stellar population of a ealaxv with those of nearby galaxies for which externally calibrated distances are available. accurate estimates of distances can be obtained up to 240 Mpc (Tonryetal.2000.2001).," Comparing SBFs produced by the stellar population of a galaxy with those of nearby galaxies for which externally calibrated distances are available, accurate estimates of distances can be obtained up to $\sim$ 40 Mpc \citep{T00,T01}."231. Other authors have used SBF studies (o determine the age and metallicity ofunresolved stellar populations of nearby galaxies (Lin.Charlot.&Graham2000:Blakeslee.Vazclekis.&Ajhar2001:Hidalgo.Marín-Franeh.Aparicio 2003): however. the SBF signal can provide information about other kinds of unudetected ancl unresolved objects in an image.," Other authors have used SBF studies to determine the age and metallicity ofunresolved stellar populations of nearby galaxies \citep{Liu00,Vaz01,HMA02}; however, the SBF signal can provide information about other kinds of undetected and unresolved objects in an image."232 This is the case. for example. of elobular cluster populations (Blakeslee&Tonrv1995:Blakeslee1999;Marín-Franeh&Aparicio2002. 2003).," This is the case, for example, of globular cluster populations \citep{BT95,B99,MA02a,MA02b}."233. In this paper. SBFs have been used to characterize the faint end of n(») using undetected galaxies in the IIDE-N images.," In this paper, SBFs have been used to characterize the faint end of $n(m)$ using undetected galaxies in the HDF-N images."234 Next in this section. the theoretical background of SBFs and the ILDE-N signal measurement are ceseribecl in detail.," Next in this section, the theoretical background of SBFs and the HDF-N signal measurement are described in detail."235 The SBF technique involves spectral analysis of the pixel-to-pixel fluctuation signal., The SBF technique involves spectral analysis of the pixel-to-pixel fluctuation signal.236 This provides the total point spread fiction PSE-convolved variance (1) produced by all point objects whose spatial {his distribution is convolved with the PSF. aud the total variance (191).," This provides the total point spread function PSF-convolved variance $P_0$ ) produced by all point objects whose spatial flux distribution is convolved with the PSF, and the total non-PSF-convolved variance $P_1$ )."237ET) A convolution in the real space (transforms into a product in Fourier space., A convolution in the real space transforms into a product in Fourier space.238 For this reason. the power spectrum of an image. P(/). has the form:," For this reason, the power spectrum of an image, $P(k)$ , has the form:"239"Further ? analyzed 22 stars using both the 923 nm triplet and the 869 nm doublet finding a flat trend from both diagnostics, but significant differences in abundances for the stars with [Fe/H]x-1.5.","Further \cite{Takada-Hidai2005} analyzed 22 stars using both the 923 nm triplet and the 869 nm doublet finding a flat trend from both diagnostics, but significant differences in abundances for the stars with $[\mathrm{Fe} / \mathrm{H}] \leq -1.5$."240" These differences might be due to non-LTE effects in the diagnostics used, and this was investigated by ? by calculating non-LTE corrections for the 869 nm doublet, the 923 nm triplet, and in addition the 1045 nm triplet."," These differences might be due to non-LTE effects in the diagnostics used, and this was investigated by \cite{Takeda2005} by calculating non-LTE corrections for the 869 nm doublet, the 923 nm triplet, and in addition the 1045 nm triplet."241 They used data for determining the evolution of non-LTE corrections to the 869 nm doublet and the 923 nm triplet (the 1045 nm triplet not being observed at the time).," They used data for determining the evolution of sulphur from \cite{Israelian2001,Takada-Hidai2002,Ryde2004,Nissen2004,Takada-Hidai2005} and some older works and applied their non-LTE corrections to the 869 nm doublet and the 923 nm triplet (the 1045 nm triplet not being observed at the time)."242" Surprisingly, they found that the discrepancies between the evolution of sulphur measured using the different diagnosticsincreased;; the 869 nm doublet suggesting a steadyincrease of [S/Fe] for lower [Fe/H] and the 923 nm triplet indicating a plateau for halo stars."," Surprisingly, they found that the discrepancies between the evolution of sulphur measured using the different diagnostics; the 869 nm doublet suggesting a steadyincrease of $[\mathrm{S} / \mathrm{Fe}]$ for lower $[\mathrm{Fe} / \mathrm{H}]$ and the 923 nm triplet indicating a plateau for halo stars."243 To resolve this mismatch and check their non-LTE modeling they proposed observations of the 1045 nm triplet., To resolve this mismatch and check their non-LTE modeling they proposed observations of the 1045 nm triplet.244" ? combined spectra from four observation runs, initially aimed at studying other elements into the largest sample of Galactic sulphur abundance measurements with a total of 74 dwarfs."," \citet{Caffau2005} combined spectra from four observation runs, initially aimed at studying other elements into the largest sample of Galactic sulphur abundance measurements with a total of 74 dwarfs."245" They used the NTT and VLT telescopes and since the wavelength coverage of the observation runs varied, they did not use the same diagnostics in all their sulphur abundance determinations."," They used the NTT and VLT telescopes and since the wavelength coverage of the observation runs varied, they did not use the same diagnostics in all their sulphur abundance determinations."246" For each star they used as many diagnostics as possible of a very weak multiplet around 675 nm, the doublet around 869 nm, and the triplet around 923 nm."," For each star they used as many diagnostics as possible of a very weak multiplet around 675 nm, the doublet around 869 nm, and the triplet around 923 nm."247 They found an [S/Fe] vs. [Fe/H] plot with a large scatter in [S/Fe] for -2.4<[Fe/H]-1 resembling a combination of the two types of evolution previously described., They found an $[\mathrm{S} / \mathrm{Fe}]$ vs. $[\mathrm{Fe} / \mathrm{H}]$ plot with a large scatter in $[\mathrm{S} / \mathrm{Fe}]$ for $-2.4 \leq [\mathrm{Fe} / \mathrm{H}] \leq -1$ resembling a combination of the two types of evolution previously described.248 The development of better infrared spectrometers has made more sulphur lines measurable and the problem with finding a suitable diagnostic for determining sulphur abundance is nowadays less severe., The development of better infrared spectrometers has made more sulphur lines measurable and the problem with finding a suitable diagnostic for determining sulphur abundance is nowadays less severe.249 One example is the 1] line at 1082 nm used in this work., One example is the ] line at 1082 nm used in this work.250" This line is not believed to be affected by non-LTE effects, but unfortunately the 1] line is undetectable in halo dwarfs for low [Fe/H]."," This line is not believed to be affected by non-LTE effects, but unfortunately the ] line is undetectable in halo dwarfs for low $[\mathrm{Fe} / \mathrm{H}]$."251 It is however detectable in giants down to [Fe/H]~—2.5., It is however detectable in giants down to $[\mathrm{Fe} / \mathrm{H}] \sim -2.5$.252 Another example is the triplet around 1045 nm also used in this work., Another example is the triplet around 1045 nm also used in this work.253" This triplet is not as strong as the triplet at 923 nm, but it is situated in a spectral region almost unaffected by telluric lines."," This triplet is not as strong as the triplet at 923 nm, but it is situated in a spectral region almost unaffected by telluric lines."254 ? investigated the possibility of using the 1045 nm triplet for sulphur determination in disk dwarfs and it was recently used by them fordetermining sulphur abundances in four halo dwarfs corroborating the scatter in [S/Fe] for [Fe/H]~—1 found in ? but at a lower level (?).., \citet{Caffau2007b} investigated the possibility of using the 1045 nm triplet for sulphur determination in disk dwarfs and it was recently used by them fordetermining sulphur abundances in four halo dwarfs corroborating the scatter in $[\mathrm{S} / \mathrm{Fe}]$ for $[\mathrm{Fe} / \mathrm{H}] \sim -1$ found in \citet{Caffau2005} but at a lower level \citep{Caffau2010}.255 ? also recently used the 1045 nm triplet to determine [S/Fe] in 33 halo/disk stars proposing yet a new scenario for the evolution of sulphur; a zig-zag trendwith a local plateau around [S/Fe]~0.3 for -2.5<[Fe/H]x-1.5 preceded by a rise in [S/Fe] for lower [Fe/H]., \cite{Takeda2010} also recently used the 1045 nm triplet to determine $[\mathrm{S} / \mathrm{Fe}]$ in 33 halo/disk stars proposing yet a new scenario for the evolution of sulphur; a zig-zag trendwith a local plateau around $[\mathrm{S} / \mathrm{Fe}] \sim 0.3$ for $-2.5 \leq [\mathrm{Fe} / \mathrm{H}] \leq -1.5$ preceded by a rise in $[\mathrm{S} / \mathrm{Fe}]$ for lower $[\mathrm{Fe} / \mathrm{H}]$.256" This is however not confirmed by ? using the 923 nm triplet to determine sulphur abundance in 33 giants and turnoff stars, with iron abundances [Fe/H]x—2.5, without finding any stars with ‘high’ [S/Fe]."," This is however not confirmed by \citet{Spite2011} using the 923 nm triplet to determine sulphur abundance in 33 giants and turnoff stars, with iron abundances $[\mathrm{Fe} / \mathrm{H}] \leq -2.5$, without finding any stars with `high' $[\mathrm{S} / \mathrm{Fe}]$."257 Obviously there is no consensus on the Galactic chemical evolution of sulphur., Obviously there is no consensus on the Galactic chemical evolution of sulphur.258 Is there a scatter or some kind of rise for halo stars or not?, Is there a scatter or some kind of rise for halo stars or not?259 In this work we will present and compare sulphur abundance measurements for 10 halo K giants using both the [51] line at 1082 nm and the 1045 nm triplet., In this work we will present and compare sulphur abundance measurements for 10 halo K giants using both the ] line at 1082 nm and the 1045 nm triplet.260 Since the [51] line is not affected by non-LTE effects we will also be able to estimate the non-LTE effects of the 1045 nm triplet by comparing the derived sulphur abundances from the two diagnostics., Since the ] line is not affected by non-LTE effects we will also be able to estimate the non-LTE effects of the 1045 nm triplet by comparing the derived sulphur abundances from the two diagnostics.261 Thereby we will be able to check the validity of the non-LTE corrections calculated by ? for halo K giants., Thereby we will be able to check the validity of the non-LTE corrections calculated by \cite{Takeda2005} for halo K giants.262 Since non-LTE corrections seem to be significant for most sulphur diagnostics this empirical test of the corrections is important., Since non-LTE corrections seem to be significant for most sulphur diagnostics this empirical test of the corrections is important.263 We have observed ten K giants in the Galactic halo using the spectrometer CRIRES (???) mounted on VLT.," We have observed ten K giants in the Galactic halo using the spectrometer CRIRES \citep{Kaufl2004,Moorwood2005,Kaufl2006} mounted on VLT."264 Giants were chosen because the 1082 nm i] line is too weak to be observed in metal-poor dwarfs (see Sect. ??))., Giants were chosen because the 1082 nm ] line is too weak to be observed in metal-poor dwarfs (see Sect. \ref{blends}) ).265 CRIRES is a high-resolution echelle spectrometer designed for near-infrared observations and it uses nodding and jittering to eliminate the sky background and adaptive optics to enhance the S/N. Basic data for the observed stars are shown in Table[I] and a summary of the observations is shown in Table []., CRIRES is a high-resolution echelle spectrometer designed for near-infrared observations and it uses nodding and jittering to eliminate the sky background and adaptive optics to enhance the S/N. Basic data for the observed stars are shown in Table \ref{tab:starinfo} and a summary of the observations is shown in Table \ref{tab:criresobs}.266 The difference in S/N in the two settings is in some cases due to different integration times and in changes due to clouds in the star's visibility between the two observations., The difference in S/N in the two settings is in some cases due to different integration times and in changes due to clouds in the star's visibility between the two observations.267" The observations were made with a slit width of 0.25"" resulting in a spectral resolution of R=A/AA~80000 and 2.5 pixels per resolution element."," The observations were made with a slit width of 0.25"" resulting in a spectral resolution of $R= \lambda / \Delta \lambda \sim 80000$ and 2.5 pixels per resolution element."268 They were carried out in service mode during October 2007 - February 2008., They were carried out in service mode during October 2007 - February 2008.269 We used two standard settings: one covering the triplet around 1045 nm and one covering the [51] line at 1082 nm., We used two standard settings: one covering the triplet around 1045 nm and one covering the ] line at 1082 nm.270 The spectral range of the two settings used are roughly 12 nm with the sulphur lines as centered as possible., The spectral range of the two settings used are roughly 12 nm with the sulphur lines as centered as possible.271 A couple of fast rotating B stars were also observed used for checking that no telluric lines were affecting the lines analyzed., A couple of fast rotating B stars were also observed used for checking that no telluric lines were affecting the lines analyzed.272" Using only a few lines, or even just one, for an abundance determination such as ours calls for an extra careful examination of blends, and in case blends are found: to try to choose the stellar sample in such a way that the effectsof the blending are minimized."," Using only a few lines, or even just one, for an abundance determination such as ours calls for an extra careful examination of blends, and in case blends are found: to try to choose the stellar sample in such a way that the effectsof the blending are minimized."273" The [O1] line at 630 nm, analogous to our [S1] line, is for example blended with a line, but the blending is not relevant for metal-poor stars (?).."," The ] line at 630 nm, analogous to our ] line, is for example blended with a line, but the blending is not relevant for metal-poor stars \citep{Allende2001}."274 To determine for which of our sulphur lines there might be significant blends we have calculated synthetic equivalent widths for all lines in the relevant wavelength region in a grid of 288 model atmospheres., To determine for which of our sulphur lines there might be significant blends we have calculated synthetic equivalent widths for all lines in the relevant wavelength region in a grid of 288 model atmospheres.275 For metals we have used a line list of a relevant wavelength section from the VALD I database (?) updatedaccording to ?.., For metals we have used a line list of a relevant wavelength section from the VALD I database \citep{Kupka1999} updatedaccording to \cite{Gustafsson2008}. .276" The lists of molecules we were provided by Bengt Edvardsson (private communication), in turn mostly compiled by Bertrand Plez, and they include CH (?),, OH (?),, CrH (?),, SiH (electronic Kuruczp), FeH (?),,"," The lists of molecules we were provided by Bengt Edvardsson (private communication), in turn mostly compiled by Bertrand Plez, and they include CH \citep{Jorgensen1996}, , OH \citep{Goldman1998}, , CrH \citep{Burrows2002}, , SiH (electronic ), FeH \citep{Dulick2003}, ,"277and obtain a very precise tomography of stellar interiors.,and obtain a very precise tomography of stellar interiors.278further assume that Ba has an RV fluctuation of o(0.dexpar) because of the Earth's turbulent atmosphere.,"further assume that $B_N$ has an RV fluctuation of $\sigma(0,\delta v_{N,ATM})$ because of the Earth's turbulent atmosphere."279" The measured RV uncertainty oe is equal to: Three examples are considered here to represeut different level of telluric contamination to the precisionnm RVT measurements: 41. n μες-ο Len Otpms-=(OeyNEMparEOUIe""AES> aud the RV measurement. is dominated by atmospheric beliaviors. this approximation applies in a wavelength region with deuse telluric line distribution: 2. if O44.«δν (Le. in very (rausparent atmosphere windows where few telluric lines exists). then Óc44=OCs5: aud the RV uncertainty is limited by stellar photon nolse: 3. for an intermediate situation. if δρ ALC Όρηv are identical. then we apply the same weight on both RV measurements. Ops7=_[s[διimas+(Ot−≻ΡΕ...διDOο...)1/214/3."," The measured RV uncertainty $\delta v$ is equal to: Three examples are considered here to represent different level of telluric contamination to the precision RV measurements: 1, if $\delta v_{rms,S}\gg\delta v_{rms,N}$, then $\delta v_{rms}={(\delta v_{N,ATM}^2+\delta280v_{rms,N}^2)^{1/2}}$, and the RV measurement is dominated by atmospheric behaviors, this approximation applies in a wavelength region with dense telluric line distribution; 2, if $\delta v_{rms,S}\ll\delta v_{rms,N}$ (i.e., in very transparent atmosphere windows where few telluric lines exists), then $\delta281v_{rms}=\delta v_{rms,S}$, and the RV uncertainty is limited by stellar photon noise; 3, for an intermediate situation, if $\delta v_{rms,S}$ and $\delta282v_{rms,N}$ are identical, then we apply the same weight on both RV measurements, $\delta v_{rms}={[\delta v_{rms,S}+{(\delta283v_{N,ATM}^2+\delta v_{rms,N}^2)^{1/2}}]/2}$."284 Iu the practical NIB. spectroscopic observations. telluric coutaiminatiou must be miuiunized iu order to reach high Doppler precision.," In the practical NIR spectroscopic observations, telluric contamination must be minimized in order to reach high Doppler precision."285 We investigate the effectiveness of two major ways to remove telluric contamination [rom tle stellar spectra with IRET: telluric line masking aud removing. and telluric line inodeliug; and removing.," We investigate the effectiveness of two major ways to remove telluric contamination from the stellar spectra with IRET: telluric line masking and removing, and telluric line modeling and removing."286 The study results are sumiuarized below., The study results are summarized below.287 'Tellurie absorption lines are not homogeneously distributed iu the NIB spectra. instead. a large number of tellurie lines are from relatively concentrated spectral regions.," Telluric absorption lines are not homogeneously distributed in the NIR spectra, instead, a large number of telluric lines are from relatively concentrated spectral regions."288 The simplest way is to uask aud remove those severely contaiuinated regious while leaving the less coutzininated stellar ines for RV measurements., The simplest way is to mask and remove those severely contaminated regions while leaving the less contaminated stellar lines for RV measurements.289 This would reduce RV measurement uncertainty caused by telluric line contamination. leadiug to improved RV precision.," This would reduce RV measurement uncertainty caused by telluric line contamination, leading to improved RV precision."290 However. if too much of the wavelength coverage 'eejon is masked aud removed. then RV uncertainty due to photon noise would increase.," However, if too much of the wavelength coverage region is masked and removed, then RV uncertainty due to photon noise would increase."291 Therefore. a balance between the uncertainty brought by telluric lines contamination and by photou noise uust be fouud and au optimal RV performance can be achievect.," Therefore, a balance between the uncertainty brought by telluric lines contamination and by photon noise must be found and an optimal RV performance can be achieved."292 We use an M9 dwarf (T4j—2100Ix. Vsinv=5kin-s| and 129) as an example to illustrate iow this maskine techuique allects the RV measurement precision with IRET.," We use an M9 dwarf $T_{\rm{eff}}$ =2400K, $V \sin{i}$ $\rm{km\cdot s}^{-1}$ and $m_J$ =9) as an example to illustrate how this masking technique affects the RV measurement precision with IRET."293 We assume a 30 uin exposure time. a waveleneth coverage [rom 800 to 1350 nm and au instrument throughput as shown in Fie. LL. ," We assume a 30 min exposure time, a wavelength coverage from 800 to 1350 nm and an instrument throughput as shown in Fig. \ref{fig:Wav_Eta}. ."294"We calculate photou-limited RV uncertainty 905,444 and 9c4; according to Equation (13))."," We calculate photon-limited RV uncertainty $\delta v_{S,rms}$ and $\delta v_{N,rms}$ according to Equation \ref{eq:overall_Doppler_2d}) )."295 We calculate Q factors. Qs aud Qa. based ou the two components. By aud Ba. in Equation (20)).," We calculate $Q$ factors, $Q_S$ and $Q_N$, based on the two components, $B_S$ and $B_N$, in Equation \ref{eq:B_atm}) )."296 The photou flux of By aud By are calculated from the inputs of the stellar type. magnitude. exposure time. instrument specifications aud telluric absorption properties.," The photon flux of $B_S$ and $B_N$ are calculated from the inputs of the stellar type, magnitude, exposure time, instrument specifications and telluric absorption properties."297 Iu practice. RV uncertainty of By is uot domiuated by photon-noise. instead. it is domiuated by atiuosphlieric behaviors such as wind. molecular column deusity change. etc.," In practice, RV uncertainty of $B_N$ is not dominated by photon-noise, instead, it is dominated by atmospheric behaviors such as wind, molecular column density change, etc."298 ? used HARPS archive data aud found that O» lines are stable toa 10 ms.| level over 6 years., \citet{Figueira2010} used HARPS archive data and found that $O_2$ lines are stable toa 10 $\rm{m\cdot s}^{-1}$ level over 6 years.299 However. the stability of," However, the stability of"300When the AGN jet moves through the ICM. it inlluences it in numerous ways.,"When the AGN jet moves through the ICM, it influences it in numerous ways."301 First. it can induce motions in large amounts of material.," First, it can induce motions in large amounts of material."302 As the LOCAL gets swept up in the shockfront. some of it moves outwards with the shocked shell ancl some σος caught in a backllow.," As the ICM gets swept up in the shockfront, some of it moves outwards with the shocked shell and some gets caught in a backflow."303 In the [ate stages of the evolution of the svstem. a significant amount of ICAL falls back into the core regions as the cocoon Collapses and buovanthy rises into the cluster atmosphere.," In the late stages of the evolution of the system, a significant amount of ICM falls back into the core regions as the cocoon collapses and buoyantly rises into the cluster atmosphere."304 All of these effects. will imprint signatures (broadening. and the development of blucshifted anc redshiftecl wings and peaks) in the absorption line profiles.," All of these effects will imprint signatures (broadening, and the development of blueshifted and redshifted wings and peaks) in the absorption line profiles."305 The second strong οσο is a change of the ICM temperature., The second strong effect is a change of the ICM temperature.306 Strong shocks associated with the carly phase of jet activity will appreciably heat some regions of the ICM. tvpically raising it to temperatures sullicient to fully ionize oxvgen.," Strong shocks associated with the early phase of jet activity will appreciably heat some regions of the ICM, typically raising it to temperatures sufficient to fully ionize oxygen."307 On the other hand. unshocked ICM that is caught in the supelralt™ froma buovant cocoon will adiabatically decompress and be cooled. (strengthening oxygen absorption features).," On the other hand, unshocked ICM that is caught in the “updraft” from a buoyant cocoon will adiabatically decompress and be cooled (strengthening oxygen absorption features)."308 Clearly. the kinematics and the thermocdvnamics. are coupled.," Clearly, the kinematics and the thermodynamics are coupled."309 We might expect (and. as discussed below. confirm) it to be appreciably harder to form high velocity oxvecn absorbers than iron absorbers. since anv strong shock responsible for accelerating the gas. will inevitably heat it lo re point where oxvecn becomes fully ionized.," We might expect (and, as discussed below, confirm) it to be appreciably harder to form high velocity oxygen absorbers than iron absorbers, since any strong shock responsible for accelerating the gas will inevitably heat it to the point where oxygen becomes fully ionized."310 On the othr hand. we find iron lines which. in the strongest cases. display a double peak structure. with one peak being at the rest energy of the line ancl one being clearly blueshifted but still attached to the line itself.," On the other hand, we find iron lines which, in the strongest cases, display a double peak structure, with one peak being at the rest energy of the line and one being clearly blueshifted but still attached to the line itself."311 In those cases. the blueshifted components arise from a [large amount of shock accelerated material oulllowing at fairly high speeds. whereas the peak at the rest energy comes from still undisturbed material in ront of the jet.," In those cases, the blueshifted components arise from a large amount of shock accelerated material outflowing at fairly high speeds, whereas the peak at the rest energy comes from still undisturbed material in front of the jet."312 Representative results for oxveen ane iron line profiles and equivalent widths are shown in Figs., Representative results for oxygen and iron line profiles and equivalent widths are shown in Figs.313 47., 4–7.314 Lt is instructive o separate the discussion of the phase of activity [rom that of the phase., It is instructive to separate the discussion of the phase of activity from that of the phase.315 During the cocoon inllation phase. he radio galaxy activity inflates an over-pressurecl Cocoon which. in turn. drives a strong shock into the ICM.," During the cocoon inflation phase, the radio galaxy activity inflates an over-pressured cocoon which, in turn, drives a strong shock into the ICM."316 The jiegh ICM temperatures produced by this shock leads to a rapid drop in the equivalent widths of the OVIL ΟΝΗΙ. SINII and FeNAY lines with rather little change in the line xofile.," The high ICM temperatures produced by this shock leads to a rapid drop in the equivalent widths of the OVII, OVIII, SiXIII and FeXXV lines with rather little change in the line profile."317 “Phe decrease in line equivalent widths are rather more dramatic for lines of sight close to the jet axis simply due to the increased. path-Iength along which the ambient inc-absorbing ICM. has been scoured out by the cocoon shock., The decrease in line equivalent widths are rather more dramatic for lines of sight close to the jet axis simply due to the increased path-length along which the ambient line-absorbing ICM has been scoured out by the cocoon shock.318 ‘Towards the end of the cocoon-inflation phase. the Cocoon pressure becomes comparable to that of the ambient ICM and the shock weakens.," Towards the end of the cocoon-inflation phase, the cocoon pressure becomes comparable to that of the ambient ICM and the shock weakens."319 After that. the source transits into the cocoon-collapse phase in which the shocked LOCAL surrounding the sides of the cocoon (ie. the equatorial regions with respect to the jet axis) falls back towards the cluster center.," After that, the source transits into the cocoon-collapse phase in which the shocked ICM surrounding the sides of the cocoon (i.e., the equatorial regions with respect to the jet axis) falls back towards the cluster center."320 Phe result of this infall is to squeeze” the cocoon plasma ancl transform. it into two buovantly rising. mushroom-shapecl plumes (see Fig.," The result of this infall is to “squeeze” the cocoon plasma and transform it into two buoyantly rising, mushroom-shaped plumes (see Fig."321" 1 of Reynolds. Ποια, Begelman 2002)."," 1 of Reynolds, Heinz Begelman 2002)."322 During these times. the equivalent widths of the OVIL OVLIE SiXIIE and. FeXXNV. absorption lines eracually recover to approximately their initial values.," During these times, the equivalent widths of the OVII, OVIII, SiXIII and FeXXV absorption lines gradually recover to approximately their initial values."323 Llowever. the complex. dvnamies of the ICM curing this phase create imprints in the line profiles.," However, the complex dynamics of the ICM during this phase create imprints in the line profiles."324 For lines of sight close to the jet axis. both the oxveen ancl iron. absorption lines develop blue wings corresponding to absorption bv the outward moving ICM shell and. latter. ICM. that is being dragged: out of the core regions of the cluster in the wake of the buovant ICM plume.," For lines of sight close to the jet axis, both the oxygen and iron absorption lines develop blue wings corresponding to absorption by the outward moving ICM shell and, latter, ICM that is being dragged out of the core regions of the cluster in the wake of the buoyant ICM plume."325 These features are subtle in the case of the oxvecn lines. but are rather dramatic in the case of the FeXXV lines.," These features are subtle in the case of the oxygen lines, but are rather dramatic in the case of the FeXXV lines."326 In all cases. the velocities characterizing the blue-wing are less than but of the order of the ICM sound speed.," In all cases, the velocities characterizing the blue-wing are less than but of the order of the ICM sound speed."327 Interestingly. for lines of sight that make a large angle with the jet-axis. the FeXXV line displavs a subtle recshift of its centroid corresponding to the actual inward. collapse of the ICM core (Fig.," Interestingly, for lines of sight that make a large angle with the jet-axis, the FeXXV line displays a subtle redshift of its centroid corresponding to the actual inward collapse of the ICM core (Fig."328 7)., 7).329 There has been much recent excitement about the possible detection. of the WIILM. through the OVILE and OVILL K-shell resonant absorption lines., There has been much recent excitement about the possible detection of the WHIM through the OVII and OVIII K-shell resonant absorption lines.330 This is of obvious importance eiven that these WIIIM. filaments are expected to be the repository for half of the barvons in the local Universe., This is of obvious importance given that these WHIM filaments are expected to be the repository for half of the baryons in the local Universe.331 The most robust detection to date was obtained by a lsh ποιον. Transmission. Cratings (LIEZEGO observation of the blazar Mrk. 421 during an outburst (Nicastro et al., The most robust detection to date was obtained by a High Energy Transmission Gratings (HETG) observation of the blazar Mrk 421 during an outburst (Nicastro et al.332 2005a.b).," 2005a,b)."333 Two absorber systems were detected in. both OVIL and NVIE at velocities of ez=3800+300kms and ez=SO90τε300 (to be compared with the recession velocity of the blazar of CZ=9000kms 5)., Two absorber systems were detected in both OVII and NVII at velocities of $cz=3300\pm 300\kmps$ and $cz=8090\pm 300$ (to be compared with the recession velocity of the blazar of $CZ=9000\kmps$ ).334 The OVILE Ka EWs of these two systems were measured to be 0.080-E0.021eV. (3.0EOc δ1) and 0.05920.021eV (2.24EOc Sm).," The OVII $\alpha$ EWs of these two systems were measured to be $0.080\pm3350.021\eV$ $3.0\pm 0.8 m\AA$ ) and $0.059\pm 0.021\eV$ $2.2\pm 0.8336m\AA$ )."337 Nicastro οἱ al. (, Nicastro et al. (3382005b) demonstrate that the column density distribution implied. by these detections is consistent. with the notion that the NLEM filaments do. indeed. balance the baryon budget of the local Universe.,"2005b) demonstrate that the column density distribution implied by these detections is consistent with the notion that the WHIM filaments do, indeed, balance the baryon budget of the local Universe."339 While this is the most robust detection of zον0 WLILLM. it was not the first.," While this is the most robust detection of $z>0$ WHIM, it was not the first."340 Fang et al. (, Fang et al. (3412002) claim a detection of the OVILE Wa line with an EW of 0.41eV. from a system with ez=16.600kms towarels the blazar PISS 2155304 (which has a systemic velocity of e;=34.800kms i). although the significance of this detection (which is based on finding a single line in a blind search. of the spectrum) has been questioned.,"2002) claim a detection of the OVIII $\alpha$ line with an EW of $0.41\eV$ from a system with $cz=16,600\kmps$ towards the blazar PKS 2155–304 (which has a systemic velocity of $cz=34,800\kmps$ ), although the significance of this detection (which is based on finding a single line in a blind search of the spectrum) has been questioned."342 Melxernan et al. (, McKernan et al. (3432003). also claim a OVILE line with a velocity of ὃς=4400kms and a EW of 2.17eV. towards the broad lino radio galaxy 3€120 (which has a systemic velocity of ez=9900kms 1).,2003) also claim a OVIII line with a velocity of $cz=4400\kmps$ and a EW of $2.17\eV$ towards the broad line radio galaxy 3C120 (which has a systemic velocity of $cz=9900\kmps$ ).344 Other recent detections of z0 WLLL Via X-ray absorption lines have been reported. towards the BL-Lac object PIs 0548322 (Barcons et al., Other recent detections of $z>0$ WHIM via X-ray absorption lines have been reported towards the BL-Lac object PKS 0548–322 (Barcons et al.345 2005) and the racio-loucl quasar 11821|643 (Mathur. Weinberg Chen 2003)," 2005) and the radio-loud quasar H1821+643 (Mathur, Weinberg Chen 2003)."346 While observational biases are evident. ib is striking," While observational biases are evident, it is striking"347is then given by eliony=(3/4).lo.,is then given by $A_{10}n_1=(3/4)A_{10}$.348 The 21 em tux of a cloud. with neutral hydrogen mass {μι at clistance D is (see e.g. Spitzer (1978)))., The 21 cm flux of a cloud with neutral hydrogen mass $M_{\rm HI}$ at distance $D$ is (see e.g. \cite{spit78}) ).349 In radio astronomy. this is quoted in terms of the line flux. f.SG)de. with SG) expressed in ly and the line-width in km |. ]S(e)do=[SG)Gefe) de. hence den (1991))) The 2lem ‘brightness ip. of an object is defined as the temperature at which a black-bocky emits the same Εαν.," In radio astronomy, this is quoted in terms of the line flux, $\int S(\nu)\,dv$, with $S(\nu)$ expressed in Jy and the line-width in km $^{-1}$ , $F=\int S(\nu)\,d\nu=\int350S(\nu)\,(\nu/c)\,dv$ , hence \cite{wakk91b}) ) The 21cm brightness $T_B$, of an object is defined as the temperature at which a black-body emits the same flux."351 The conversion from [lux to brightness temperature is then given by 7g/9S=4. where the telescope-dependent conversion. factor we use is R=O.158h in order to match the observational sUrvery we are comparingJv our results against Wakker (198822).," The conversion from flux to brightness temperature is then given by $T_B/S=R$, where the telescope-dependent conversion factor we use is $R=0.158 {\rm K\,Jy}^{-1}$ in order to match the observational survery we are comparing our results against \cite{huls88}) )."352 To compute the simulated. [flux we place an observer in the simulated. galaxy and evaluate the net Hux received. by an ideal radio telescope with a beam size 6., To compute the simulated flux we place an observer in the simulated galaxy and evaluate the net flux received by an ideal radio telescope with a beam size $\theta$.353 For a single SPII particle at. position r;. the fraction dimm of mass that falls within the beam at distance between r and rk|dr is where fis the smoothing length of the particle. and VM the SPL kernel.," For a single SPH particle at position ${\bf r}_i$, the fraction $dm/m$ of mass that falls within the beam at distance between $r$ and $r+dr$ is where $h$ is the smoothing length of the particle, and $W$ the SPH kernel."354 Phe total lux received from this particle is computed from I5q. (1)).," The total flux received from this particle is computed from Eq. \ref{eq:flux}) ),"355" and is represented by a Ciaussian emission line centred at velocity νιr/r with width 0,=(ApTim, 7."," and is represented by a Gaussian emission line centred at velocity ${\bf v}\cdot{\bf r}/r$ with width $\sigma_v=(k_{\rm B}\,T/m_h)^{1/2}$ ."356 Phe total spectrum is obtained integrating over dr and summing over all particles., The total spectrum is obtained integrating over $dr$ and summing over all particles.357 Simulations have been performed at two clilferent mass resolutions to ensure that resolution cllects do not allect our results., Simulations have been performed at two different mass resolutions to ensure that resolution effects do not affect our results.358 Observations have been repeated for a number of observers along the solar circle. and at times spanning a period of Ον," Observations have been repeated for a number of observers along the solar circle, and at times spanning a period of 1Gyr."359 We use these observers to compute error bars on the mock observations., We use these observers to compute error bars on the mock observations.360 The distribution of at galactic latitudes. b. greater than 20° does not depend stronglv on the time at which observations are mace. showing that the galaxy our results do not represent a transient feature.," The distribution of at galactic latitudes, $b$, greater than $^\circ$ does not depend strongly on the time at which observations are made, showing that the galaxy our results do not represent a transient feature."361 Within the galactic disk (6< 20°) the mean brightness temperature of the eas decreases slowly over time as the σας disk is converted into stars., Within the galactic disk $b<20^\circ$ ) the mean brightness temperature of the gas decreases slowly over time as the gas disk is converted into stars.362 The all-sky 21 em brightness distribution of the simulated ealaxy looks remarkably similar to the observed iin the Milkv. Way. as measured by the onn (LAL. Ixalberlaatal (2005))) survey (Fig. 1)).," The all-sky 21 cm brightness distribution of the simulated galaxy looks remarkably similar to the observed in the Milky Way, as measured by the Leiden-Argentine-Bonn (LAB, \cite{kalb05}) ) survey (Fig. \ref{fig:allsky}) )."363 Phe LAB survey has angular resolution of 0.50.57 but unfortunately our numerical simulation does not have enough particles to resolve structures on such small scales., The LAB survey has angular resolution of $0.5^\circ \times 0.5^\circ$ but unfortunately our numerical simulation does not have enough particles to resolve structures on such small scales.364 We calculate the mean angular extent of the particles that contribute to the Dux in a given direction on the sky in the simulated galaxy. and then smooth the LAB survey with a Gaussian kernel of the same width.," We calculate the mean angular extent of the particles that contribute to the flux in a given direction on the sky in the simulated galaxy, and then smooth the LAB survey with a Gaussian kernel of the same width."365 At low galactic latitudes this smoothing angle is typically less than VY., At low galactic latitudes this smoothing angle is typically less than $1^\circ$.366 Llowever. at high latitudes the mean smoothing length is much larger. and at |b]>60 can reach up to 207 due to the relatively small number of SPL particles at these Latitucdoes.," However, at high latitudes the mean smoothing length is much larger, and at $|b|>60^\circ$ can reach up to $20^\circ$ due to the relatively small number of SPH particles at these latitudes."367 Both observed ancl simulated brightness maps display a bright and thin delisk in the plane of the MW. embedded in a thicker cooler envelope (Le~10? IK). with an even cooler componcnt (Pp100 IW) at high velocities. Uu]27100 km s+ with respect to the local standard of rest (LSI).," Both observed and simulated brightness maps display a bright and thin disk in the plane of the MW, embedded in a thicker cooler envelope $T_B\sim 10^3$ K), with an even cooler component $T_B\sim368100$ K) at high velocities, $|v_{\rm lsr}| > 100$ km $^{-1}$, with respect to the local standard of rest (LSR)."369 The brightness temperature Zg. and its fall-olf with latitude. is very similar in the observed ancl simulated maps.," The brightness temperature $T_B$, and its fall-off with latitude, is very similar in the observed and simulated maps."370 The minimum brightness temperature in the simulated. map is 95lx. in &ood agreement with observations of the MW. where it is found that every linc-of-sight. contains casily observableHL.," The minimum brightness temperature in the simulated map is 95K, in good agreement with observations of the MW, where it is found that every line-of-sight contains easily observable."371 Phe simulated high velocity gas Ονc100 km s1) forms a nearly uniform background., The simulated high velocity gas $|v_{{\rm lsr}}|>100$ km $^{-1}$ ) forms a nearly uniform background.372 Due to the relatively small number of particles that are Hageged as high. velocity at any one time the spatial resolution is poor. especially at high latitudes. and these simulations do not resolve the fine structure seen in the LAB.," Due to the relatively small number of particles that are flagged as high velocity at any one time the spatial resolution is poor, especially at high latitudes, and these simulations do not resolve the fine structure seen in the LAB."373 When the LAB cata are smoothed to the same resolution as the simulation the resulting clistribution matches closely. with a mean brightness temperature. of Tp~23.22 Ix. as compared to 20.33 Ix. in the simulations.," When the LAB data are smoothed to the same resolution as the simulation the resulting distribution matches closely, with a mean brightness temperature of $T_B\sim 23.22$ K, as compared to 20.33 K in the simulations."374 The velocity cistribution of the simulated also matches well with the LAB data (Fig 2))., The velocity distribution of the simulated also matches well with the LAB data (Fig \ref{fig:veldist}) ).375 Although we cannot resolve individual HVC's. the simulated velocities are in good agreement with the properties of the ΗΝ catalogue of Lockmanetal (2002)..," Although we cannot resolve individual HVCs, the simulated velocities are in good agreement with the properties of the HVC catalogue of \cite{lock02}. ."376 Lockmanetal(2002). identified some of the detections in this survey with external galaxies. we have removed these from the plot.," \cite{lock02} identified some of the detections in this survey with external galaxies, we have removed these from the plot."377 Additionally. clouds that were identified as being part of the Alagellanic stream. which dominate the extreme negative velocity [low (Alathewson.Cleary.&Alurray (1974))). were LPOCDLOVCG.," Additionally, clouds that were identified as being part of the Magellanic stream, which dominate the extreme negative velocity flow \cite{math74}) ), were removed."378 The line flux. f.SQv)dv. for high velocity gas. 100kms 1 is shown in Fig. 3..," The line flux, $\int S(\nu)\,d\nu$, for high velocity gas, $|v_{{\rm379lsr}}|>100 {\rm \,km\,s}^{-1}$ , is shown in Fig. \ref{fig:hist}."380 Solid black lines represent the results of Wakker(1991) with the emission due to the Magellanic stream and outer arm of the MW. removed. as we do not expect an isolated galaxy to match these features.," Solid black lines represent the results of \cite{wakk91} with the emission due to the Magellanic stream and outer arm of the MW removed, as we do not expect an isolated galaxy to match these features."381 Errors on the mock data are caleulated. by repeating the measurements for observers at. different points along the solar circle., Errors on the mock data are calculated by repeating the measurements for observers at different points along the solar circle.382 Phe mock and real cata look remarkably similar., The mock and real data look remarkably similar.383 The distribution of neutral gas perpendicular to the galactic plane is approximately exponential with a scale height. of 5 kpc. in agreement with the predictions of Bregman(1980) for a galactic fountain.," The distribution of neutral gas perpendicular to the galactic plane is approximately exponential with a scale height of 5 kpc, in agreement with the predictions of \cite{breg80} for a galactic fountain."384 Although as noted in Fig 3. the distribution of 21cm emission in galactic latitude is slightly more concentrated in the simulated data than the observed data. when averaged over the whole skv the mean values of the smoothed maps agree to within and. are901 for the observed data. and 4461Ix. for the simulated data.," Although as noted in Fig \ref{fig:hist} the distribution of 21cm emission in galactic latitude is slightly more concentrated in the simulated data than the observed data, when averaged over the whole sky the mean values of the smoothed maps agree to within and are485K for the observed data, and 446K for the simulated data."385 lt is οΠο to measure the distances to HIVC's in the real universe., It is difficult to measure the distances to HVCs in the real universe.386 However. in our simulations this information is preserved and can be measured easily (Fig 4)).," However, in our simulations this information is preserved and can be measured easily (Fig \ref{fig:dist}) )."387 of the Hus is emittedwithin a distance of 13 kpe from our observer., of the flux is emittedwithin a distance of 13 kpc from our observer.388index pairs containing at least one ( were obtained from CLTHfOBIVEVEReEB.MVUB de}.,"index pairs containing at least one $U$ were obtained from ${\rm CI}_i\vert_{i=1}^{i=10}=389\{U-2B+V,U-V,U-R_C,U-I_C,B-V,B-R_C,B-I_C,V-R_C,V-I_C,R_C-I_C\}$ ."390 Taking Ad]=1.60. £(BVV)=0.015 (Liu&Janes1990). the basic quantities (10)) are summarized in Table 1. for different: phases.," Taking $[M]=-1.60$, $E(B-V)=0.015$ \citep{liuj1} the basic quantities for \ref{107a}) ) are summarized in Table \ref{tab1} for different phases."391 lig. lis a plot of Πίο). 962). loggel). ο) for one whole pulsation.," Fig.1 is a plot of $T_{\rm e}(\varphi)$, $\vartheta(\varphi)$, $\log g_{\rm e}(\varphi)$, $h_0(R,\varphi)$ for one whole pulsation."392 ελ) and Alogq.(s) are plotted in the lower panels of Fig. 2, $\Delta T_{\rm e}(\varphi)$ and $\Delta\log g_{\rm e}(\varphi)$ are plotted in the lower panels of Fig. \ref{fig2}.393 ssumineg random errors of £0.02 for the colour indices will result in AY.=Εθν and Aloeg.=+£0.04., Assuming random errors of $\pm 0.02$ for the colour indices will result in $\Delta T_{\rm e}=\pm 10\mbox{K}$ and $\Delta\log g_{\rm e}=\pm 0.04$.394 These values are indicated by the dotted horizontal lines., These values are indicated by the dotted horizontal lines.395 Condition Lis satisfied in the phase points Lvine below or close to the clotted lines., Condition I is satisfied in the phase points lying below or close to the dotted lines.396 At phases [ving. above the dotted lines. the monochromatic lux of static models and SU Dra dillers significantly on a level which has noticeable ellect on the broad band colours (BV (De: ," At phases lying above the dotted lines, the monochromatic flux of static models and SU Dra differs significantly on a level which has noticeable effect on the broad band colours $UBV(RI)_C$ ."397The assumed Al] L6.E(BVW)=0.015 were verified by a variation procedure (Bareza&Benkó2009).," The assumed $[M]=-1.6$, $E(B-V)=0.015$ were verified by a variation procedure \citep{barc3}."398. In the shock free phases yo=0.15.0.5.0.55. minimization of AT.(uz).Adoggol) resulted in A]= 100-010. E(BV)=0.015£0.01.," In the shock free phases $\varphi=0.15,0.5,0.55$, minimization of $\Delta399T_{\rm e}(\varphi),\Delta\log g_{\rm e}(\varphi)$ resulted in $[M]=-1.60\pm 0.10$ , $E(B-V)=0.015\pm 0.01$."400 To demonstrate the difference between good ancl poor QSAA. the histograms of the 30 logg..7; values are plotted in the upper panels of Fig.," To demonstrate the difference between good and poor QSAA, the histograms of the 30 $\log g_{\rm e},T_{\rm e}$ values are plotted in the upper panels of Fig."401 2 for 4;=0.5.0.98 of SU Dra and BD |67 τος.," \ref{fig2} for $\varphi=0.5,0.98$ of SU Dra and BD +67 708."402 Phev show normal distributions with small scatter for the non-variable BD 708 and SU Dra at y=0.5., They show normal distributions with small scatter for the non-variable BD 708 and SU Dra at $\varphi=0.5$.403 At y=0.98. the distribution is almost uniform.," At $\varphi=0.98$, the distribution is almost uniform."404 At the next phase point ~=1. merely 14 intersections of CES(loggo.Chi.Clo.AL.E(BΕεντος be. 14. pairs of loggo.T; were found instead of 30 pairs.," At the next phase point $\varphi=1$, merely 14 intersections of $\{T_{\rm e}^{(i)}(\log g_{\rm e},{\rm CI}_1,{\rm CI}_2,[M],E(B-V))\}_405{i=1,2}$, i.e. 14 pairs of $\log g_{\rm e},T_{\rm e}$ were found instead of 30 pairs."406 Phus. at 4;21 the observed. colours diller significantly from those of any static model of Ixurucz(1997).," Thus, at $\varphi \approx 1$ the observed colours differ significantly from those of any static model of \citet{kuru1}."407. Llowever. it is interesting to note that the small errors Aloeg.=0.03. AT.=191 indicate ⋜↧↓≻↓↥⋜↧≻∢⊾↕≻↓⋜↧↓↥∠⇂⋜↧↥↴∣↘⊤∶∪⋅≤⋗∶∫≻∖∖⋰⊔↓↥∪∣⋡≱∖⋖⋅↓⋅∖⇁∢⋅∠⇂⋜⋃⊔⇂⊳∖↿⋜∐⊀⊔⇍⊔↓⋯⇂⋖⋅⇂ ≼↛∪↓∪⊔↓⋅⊳∖⊲↓⊔⋜↧⋏∙≟↓⋅∢⊾∢⊾⊔↓∢⊾⊔⇂⇂⋅∪↓⋅⊳∖∪⊔↓⋖⋅∪⋅∪∐↗≈↓∪⊔↓⊲↓⊔⋯∢⊾⊳∖⊳↾↓∖↓↕⊀↓⊳∖ ↓⋯↓≻↓≻∢⊾," However, it is interesting to note that the small errors $\Delta\log g_{\rm e}=0.03$, $\Delta T_{\rm e}=12\mbox{K}$ indicate a phase island at $\varphi=0.93$ with observed and static model colours in agreement for some $0.01P\approx 10$ minutes."408⊔≱∖↥∪∣⋈⋅↥⇂↥⋖⋅↓≻∐∥⊳∖∢⊾∪⇂⋅↿↓∐⋅↓⋯⊔↓↓≻∪⊔↿↓↕∢⋅↓⊀↓⋏∙≟↓∐≼∼⊔↓⋅∖⇁∢⊾∖∖⋎↓↕⋖⋅⊔ the inward and outward motions encounter and. produce a shock (Smith1995).., This happens to be the phase of the hump on the light curve when the inward and outward motions encounter and produce a shock \citep{smit1}.409 In the interval 0.92«45<1.05 the atmosphere is in à state of maximal compression by the shock coming from the sub-photospheric lavers and {1 is nearly minimal., In the interval $0.92 < \varphi < 1.05$ the atmosphere is in a state of maximal compression by the shock coming from the sub-photospheric layers and $R$ is nearly minimal.410 ‘This is the risine branch and the start of the descending branch in the light curve., This is the rising branch and the start of the descending branch in the light curve.411 Therefore. the values of logg..ἐν. obtained from QSSA. if they can be found at all. must be considered as a first approximation only.," Therefore, the values of $\log g_{\rm e},T_{\rm e},\vartheta$ obtained from QSSA, if they can be found at all, must be considered as a first approximation only."412 This is rellected in large Aloe q«. AZ. except for qz0.93.," This is reflected in large $\Delta \log g_{\rm e}$ , $\Delta T_{\rm e}$ except for $\varphi \approx 0.93$."413 To obtain ο. £). ete.," To obtain $v(r,t)$ , etc.,"414 For (10)). #0) and PoGR.1) were differentiated by midpoint. formulae.," for \ref{107a}) ), $\vartheta(t)$ and $h_0(R,t)$ were differentiated by midpoint formulae."415 Fig. 3((, Fig. \ref{fig3}( (416a) is à. plot of thefunctions ία. for the phases /=P. s=0.15-0.35.0.5.0.55 with E).a7(R4)=0.,"a) is a plot of thefunctions ${\cal M}(d,t)$ for the phases $t=\varphi P$, $\varphi=0.15\mbox{-}0.35,0.5,0.55$ with $a^{\rm (dyn)}(R,t)=0$."417" Lhe average and stancard error of .M.d are given in Table 2. (rom the pairs and ge=0.15.0.3.0.35.0.55). (4,=0.35 and ) in (10)) as our best. values denoted by "," The average and standard error of ${\cal M},d$ are given in Table \ref{tab2} from the pairs $(\varphi_1=0.25$ and $\varphi_2=0.15,0.3,0.35,0.55)$, $(\varphi_1=0.35$ and $\varphi_2=0.55)$ in \ref{107a}) ) as our best values denoted by $[\ast]$."418At qj0.35 Condition Lis moderately violated. but Condition LL is satisfied and a(2.1D)(1.0x0.07: therefore. this phase was included to obtain dM of x].," At $\varphi=0.35$ Condition I is moderately violated, but Condition II is satisfied and $a^{\rm (dyn)}(R,t)/g_{\rm419 s}(R,t)\approx -0.07$; therefore, this phase was included to obtain $d,{\cal M}$ of $[\ast]$ ."420 Condition I is satisfied at ο=0.2.0.5: however. these phases had to be excluded from themass and distance determination because of the Large aP(gu=(2.05)25.8.A3ms7 and (q—0 00.52.38. respectively.," Condition I is satisfied at $\varphi=0.2,0.5$; however, these phases had to be excluded from themass and distance determination because of the large $a^{\rm421 (dyn)}(R,\varphi=0.2,0.5)=5.8,-3.3\:\mbox{ms}^{-2}$ and $q=-0.52,38$ , respectively."422 Using our best solution for .Vf and d. the radius variation. velocities and the components of acceleration werecomputed in physical units and are plotted in Pig. 3((," Using our best solution for ${\cal M}$ and $d$ , the radius variation, velocities and the components of acceleration werecomputed in physical units and are plotted in Fig. \ref{fig3}( ("423b)-(d).,b)-(d).424 Velocities and accelerations are plotted only for the phases of more or less good QSAA (4=0.15-0.3.0.5. 0.55). including theslightly shockedphases y= 0.35-0.45.," Velocities and accelerations are plotted only for the phases of more or less good QSAA $\varphi=0.15\mbox{-}0.3,0.5,0.55$ ), including theslightly shockedphases $\varphi=0.35\mbox{-}0.45$ ."425 At the phases v0.10.0.95.0.9.090.0.55 aPRu)=TOS.16.65.93ems7 and Condition Lis satisfied. aPR.fgGID|< 0.13. Le.AMORSs) is small in comparisonwith the other acceleration. ternis in (3).," At the phases $\varphi=0.15,0.25,0.3,0.35,0.55$ $a^{\rm (dyn)}(R,\varphi)=-79,8,-16,65,93\:\mbox{cms}^{-2}$ and ConditionI is satisfied, $\vert a^{\rm (dyn)}(R,t)/g_{\rm s}(R,t)\vert < 0.13$ , i.e.$a^{\rm (dyn)}(R,\varphi)$ is small in comparisonwith the other acceleration terms in \ref{1.100}))."426 |g)&O.L is expected. from Alogg.= 0.04.," $\vert q\vert\approx 0.1$ is expected from $\Delta \log g_{\rm e}=0.04$ ,"427regions: in cach case the Lree-Lree emission is dominated byIli.,; in each case the free-free emission is dominated by.428 Collisions between a free electron ancl trace aatom sometimes excite the aatom. which then radiates away the excitation energy.," Collisions between a free electron and trace atom sometimes excite the atom, which then radiates away the excitation energy."429 The collisional excitation rate coellicicnt for transitions from the n=l] state to state m. quia. ds where Tis in Ix. O(1.0) is the (temperature-cepencdent) elective collision strength for transitions from the η=1 to state n. oy is the statistical weight of the 7=1 state and vL0) is the energy dilference between the ài—1 and sta n (Osterbrock 1989)..," The collisional excitation rate coefficient for transitions from the $n=1$ state to state $u$, $q_{1u}$ , is where $T$ is in $\kelvin$, $\Omega(1,u)$ is the (temperature-dependent) effective collision strength for transitions from the $n=1$ to state $u$, $\omega_1$ is the statistical weight of the $n=1$ state and $\chi(1,u)$ is the energy difference between the $n=1$ and state $u$ \cite{ost89}. ."430 We compute the collisional excitation to the m=2 and n—S states using cross-sections [from Callaway (1985). ancl Callaway. Unnikrishnan Oza (1987).," We compute the collisional excitation to the $n=2$ and $n=3$ states using cross-sections from Callaway \shortcite{cal85} and Callaway, Unnikrishnan Oza \shortcite{cal87}."431. These excitations resultin additional Lye emission: excitations to the 2p. 3s and 3d states racdiatively decay ton=1 via Lya photons. and atoms in the 2s state (resulting from collisions to either the 2s or 3p state) may be additionally collisionally excited to the 2p state.," These excitations resultin additional $\lya$ emission: excitations to the $2p$ , $3s$ and $3d$ states radiatively decay to $n=1$ via $\lya$ photons, and atoms in the $2s$ state (resulting from collisions to either the $2s$ or $3p$ state) may be additionally collisionally excited to the $2p$ state."432 The rate of collisional Lye emission per unit stellar mass. gp. vell15 where and The factor of 0.33 accounts for 2s to 2p collisional excitation (OsterbrockLOSO).," The rate of collisional $\lya$ emission per unit stellar mass, $q_\lya^{\mathrm coll}$, is where and The factor of 0.33 accounts for $2s$ to $2p$ collisional excitation \cite{ost89}."433". Summing over the aanel rregions.. qj,coll—54510wous1 M, "," Summing over the and regions, $q_\lya^{\mathrm434coll}=3.4\times10^{47}~\se^{-1}\,\mathrm{M^{-1}_{\odot}}$ ."435lor the fi.=0 case. the total specific Luminosity per unit stellar mass emitted from a sstar ancl nebula and scattered in the LGAL (2(2). has three components that we treat: the truncated. stellar spectrum. U. the free-free spectrum. (fo and the scattered. Lya spectrum. flue(i: The truncated stellar spectrum is The £7 spectrum ds slightly. mocified between 912 and 1216A bv scattering in the IGM. when the photons in that range are cosmologically recdshiltecl into the Lya resonance (e.g. Peebles 1993).," For the $f_{\mathrm esc}=0$ case, the total specific luminosity per unit stellar mass emitted from a star and nebula and scattered in the IGM, $l_{\nu}^0(z)$, has three components that we treat: the truncated stellar spectrum, $l_{\nu}^{*{\mathrm a}}$, the free-free spectrum, $l_{\nu}^{{\mathrm ff}}$, and the scattered $\lya$ spectrum, $l_{\nu}^{\lya}(z)$: The truncated stellar spectrum is The $l_{\nu}^{*{\mathrm a}}$ spectrum is slightly modified between $912$ and $1216~\ang$ by scattering in the IGM, when the photons in that range are cosmologically redshifted into the $\lya$ resonance (e.g., Peebles 1993)."436 This cllect is small when Lifes)SO22510HHz. bes the width of the scattered Lye line (Loeb&νο1999) is small compared. to the frequency dillerence between 912A and 1216A.," This effect is small when $1.75\,\nu_{*}(z)\ll8.22\times10^{14}~\hz$, i.e., the width of the scattered $\lya$ line \cite{loe99} is small compared to the frequency difference between $912~\ang$ and $1216~\ang$."437" Since L75v.f2=30)s.2710""13Iz. we ignore. this. correction. to the spectrum."," Since $1.75\,\nu_{*}(z=30)=8.27\times10^{13}~\hz$ , we ignore this correction to the spectrum."438 The free-free spectrum is given by eq. (19))., The free-free spectrum is given by eq. \ref{ffspeceq}) ).439 The Lye spectrum is where diya is the rate of Lye photons produced. per solar mass of the ionizing star., The $\lya$ spectrum is where $q_\lya$ is the rate of $\lya$ photons produced per solar mass of the ionizing star.440 Phat rate is where du;= Qgi/M. and quí;=Quo/Ale.," That rate is where $q_{\mathrm HI}\equiv Q_{\mathrm HI}/M_*$ and $q_{\mathrm441HeII}\equiv Q_{\mathrm HeII}/M_*$."442 The factor 0.75. represents. the fraction of hydrogen recombinations that result in Lye photons. and the factor of 1.7 accounts for the number of hydrogen ionization per rrecombination. computed above.," The factor $0.75$ represents the fraction of hydrogen recombinations that result in $\lya$ photons, and the factor of $1.7$ accounts for the number of hydrogen ionization per recombination, computed above."443 show /2. ήMT (z) and ος}; for z=15.," show $l_{\nu}^{*{\mathrm a}}$, $l_{\nu}^{{\mathrm ff}}$, $l_{\nu}^{\lya}(z)$ and $l_{\nu}^0(z)$ for $z=15$."444 In this model no iionizing photons escape to the IGM. thus sstars don't contribute to reionization of the Universe (cf.," In this model no ionizing photons escape to the IGM, thus stars don't contribute to reionization of the Universe (cf."445 section 3.3)), section \ref{allescsec}) ).446 The second case we consider is that the nebula plays no role in reprocessing ionizing radiation from a sstar. feas=1.," The second case we consider is that the nebula plays no role in reprocessing ionizing radiation from a star, $f_{\mathrm447esc}=1$."448 This may be because. in contrast to the assumptions we mace in 83.2.1.. the density of the nebula is low. or because the nebula is eclumped into high-density regions with a small covering fraction. or because the nebula was blown zwav by the star(s).," This may be because, in contrast to the assumptions we made in \ref{nebpropsec}, the density of the nebula is low, or because the nebula is clumped into high-density regions with a small covering fraction, or because the nebula was blown away by the star(s)."449 Because the timescales of important ICM. processes extend. bevond the lifetime ofa sstar. in this section we will name the redshift of formation ofa sstar z;. and then describe the evolution as a function of 2.," Because the timescales of important IGM processes extend beyond the lifetime ofa star, in this section we will name the redshift of formation of a star $z_{\mathrm i}$, and then describe the evolution as a function of $z$."450 We assume the ICM is uniform with barvon censity meals)=LT.lO(1|eyem (which ignores the small fraction of barvons in collapsed. haloes). Yo=0.75 and Y=0.25.," We assume the IGM is uniform with baryon density $n_{\mathrm IGM}(z)=4511.7\times10^{-7}(1+z)^3~\cmt$ (which ignores the small fraction of baryons in collapsed haloes), $X=0.75$ and $Y=0.25$."452 lonizing photons from a sstar stream into the LGOAL and form an ionized region., Ionizing photons from a star stream into the IGM and form an ionized region.453 Because the densityis low. for z;X;30 recombinations are of little importance on the timescale of the star's lifetime. το2510° ve.," Because the densityis low, for $z_{\mathrm i}\la30$ recombinations are of little importance on the timescale of the star's lifetime, $\tau \simeq 2\times10^6~\yr$ ."454 Por thepurposes of calculating the properties of the ionized regionof the ICM. we assume that all of the ionizations occur immecdiatelv: at worst this contributes less than a3 per cent error to the computed spectrum.," For thepurposes of calculating the properties of the ionized regionof the IGM, we assume that all of the ionizations occur immediately; at worst this contributes less than a3 per cent error to the computed spectrum."455 rrecombines more quickly than Π.. ancl rrecombinations occur on a comparabletimescale to those," recombines more quickly than , and recombinations occur on a comparabletimescale to those"456"We observed M82 with the VLA on 2009 April 27 at 1.4. 4.8. 8.4, 22. and 43 GHz.","We observed M82 with the VLA on 2009 April 27 at 1.4, 4.8, 8.4, 22, and 43 GHz."457 The total observing time was 4 hours., The total observing time was 4 hours.458 We observed with two frequeney bands of 50 MHz. each in dual eircular polarization.," We observed with two frequency bands of 50 MHz, each in dual circular polarization."459 448 and 31048-7143 were used as primary flux density and phase calibrators. respectively.," 48 and J1048+7143 were used as primary flux density and phase calibrators, respectively."460 At 1.4. 4.8. and 8.4 GHz. we used a switching cycle of six minutes. spending one minute on the phase calibrator and five minutes on M82.," At 1.4, 4.8, and 8.4 GHz, we used a switching cycle of six minutes, spending one minute on the phase calibrator and five minutes on M82."461 We repeated these cycles 5 times over the observations. yielding an integration time of ~25 minutes at each frequeney.," We repeated these cycles 5 times over the observations, yielding an integration time of $\sim$ 25 minutes at each frequency."462 At 22 and 43 GHz. we used a switching cycle of three minutes. spending one minute on the phase calibrator and two minutes on M82.," At 22 and 43 GHz, we used a switching cycle of three minutes, spending one minute on the phase calibrator and two minutes on M82."463 These cycles were repeated 10 times during the observation. yielding an integration time of ~20 minutes at both frequencies.," These cycles were repeated 10 times during the observation, yielding an integration time of $\sim$ 20 minutes at both frequencies."464 The data reduction was performed in AIPS and involved amplitude calibration of 448 using source models., The data reduction was performed in AIPS and involved amplitude calibration of 48 using source models.465 Then we calibrated the phases using J1]048+7143 and made one phase and amplitude self-calibration on J1048+7143., Then we calibrated the phases using J1048+7143 and made one phase and amplitude self-calibration on J1048+7143.466 The calibration was then transferred to the target source M82., The calibration was then transferred to the target source M82.467 We performed one phase self-calibration on M82 at all frequencies except 43 GHz. where the source is too weak.," We performed one phase self-calibration on M82 at all frequencies except 43 GHz, where the source is too weak."468 The Chandra X-ray observatory (CXO) observed M82 on 2008 October 4. and 2009 April 17 and 29.," The Chandra X-ray observatory (CXO) observed M82 on 2008 October 4, and 2009 April 17 and 29."469 Each observation was taken with an exposure of about 18 ks., Each observation was taken with an exposure of about 18 ks.470 In these observations. the target is off the optical axis by more than 3 aremin. where the point spread function looks like an extended ellipse covering multiple pixels.," In these observations, the target is off the optical axis by more than 3 arcmin, where the point spread function looks like an extended ellipse covering multiple pixels."471 M82 was also observed with the Swift X-ray Telescope (XRT:??) on 2007 January 26. 2008 May |. and 2009 April 25 with exposures of 4.6 ks. 5.0 ks. and 4.7 ks. respectively.," M82 was also observed with the Swift X-ray Telescope \citep[XRT;][]{GehrelsChincariniGiommi2004,BurrowsHillNousek2005}472 on 2007 January 26, 2008 May 1, and 2009 April 25 with exposures of 4.6 ks, 5.0 ks, and 4.7 ks, respectively."473 Due to the low resolution of the telescope. the collection of known X-ray sources looks point-like in the XRT data. so count rates from the conglomerate of sources were measured.," Due to the low resolution of the telescope, the collection of known X-ray sources looks point-like in the XRT data, so count rates from the conglomerate of sources were measured."474 The data suffer from heavy pile-up. which i5 accounted for by extracting count rates from an annular region around the central. piled-up location7.," The data suffer from heavy pile-up, which is accounted for by extracting count rates from an annular region around the central, piled-up location."475. We imaged the data of 220081z from 2008 May 03 with a slightly super-resolved beam of 0.2 » 0.2 mas., We imaged the data of 2008iz from 2008 May 03 with a slightly super-resolved beam of 0.2 $\times$ 0.2 mas.476 The peak flux density was 5.9 mJy beam. the total emission was 32 mJy and we achieved an image rms of 79 jy beam!.," The peak flux density was 5.9 mJy $^{-1}$, the total emission was 32 mJy and we achieved an image rms of 79 $\mu$ Jy $^{-1}$."477 The data from 2009 April 08 were imaged with naturalweighting and a circular beam of 0.5 x 0.5 mas to have more sensitivity for extended emission., The data from 2009 April 08 were imaged with naturalweighting and a circular beam of 0.5 $\times$ 0.5 mas to have more sensitivity for extended emission.478 Here. the peak flux density was 0.36 mJy . the total emission was 4.3 mJy. and we achieved an image rms of41 py beam|.," Here, the peak flux density was 0.36 mJy $^{-1}$, the total emission was 4.3 mJy, and we achieved an image rms of 41 $\mu$ Jy $^{-1}$ ."479 The supernova was clearly detected in both epochs (see Fig. 1)., The supernova was clearly detected in both epochs (see Fig. \ref{fig:hsa}) ).480 The source is already clearly resolved in. the observation on 2008 May 03 and shows a ring like structure. typical for a radio supernova.," The source is already clearly resolved in the observation on 2008 May 03 and shows a ring like structure, typical for a radio supernova."481 In the following eleven months the source expanded and faded significantly., In the following eleven months the source expanded and faded significantly.482 M$2 was imaged at all frequencies using only data from baselines larger than 30 kt., M82 was imaged at all frequencies using only data from baselines larger than 30 $\lambda$ .483 This ensures that most ofthe, This ensures that most ofthe484low angular momentum. is spatially extended far out into the halo. and is generally nearly as old as the universe itself),"low angular momentum, is spatially extended far out into the halo, and is generally nearly as old as the universe itself."485 The red group is relatively metal rich. shows rapid rotation at all raclii. aud possesses a disk like spatial distribution wilh a concentration towards the Galactic center.," The red group is relatively metal rich, shows rapid rotation at all radii, and possesses a disk like spatial distribution with a concentration towards the Galactic center."486 The diversity between (he (wo groups suggests separate stages of formation and in our chemical evolution model each group is evolved separately., The diversity between the two groups suggests separate stages of formation and in our chemical evolution model each group is evolved separately.487 In the following model. the main thrust is consideration of the effects of globular cluster formation on chemical evolution aud not the formation mechanism itself.," In the following model, the main thrust is consideration of the effects of globular cluster formation on chemical evolution and not the formation mechanism itself."488 For the latter subject we reler to the more sophisticated modelling of Cote et ((2000.2002). Beaslev et ((2002). and Ixravtsov Gnedin (2003).," For the latter subject we refer to the more sophisticated modelling of Cote et (2000,2002), Beasley et (2002), and Kravtsov Gnedin (2003)."489 Like (he above cited works. ours is a bottom-up picture with (he formation mechanism most resembling that for the metal rich clusters in the Beasley οἱ al.," Like the above cited works, ours is a bottom-up picture with the formation mechanism most resembling that for the metal rich clusters in the Beasley et al."490 work (i.e. collisions between eas rich clumps)., work (i.e. collisions between gas rich clumps).491 For simplicity we assume (hat all clusters are formed in this way., For simplicity we assume that all clusters are formed in this way.492 In order to account [ον the bimodality. we argue Chat the collapse was anisotropic.," In order to account for the bimodality, we argue that the collapse was anisotropic."493 The objective of this work is to reconstruct the cosmic star Formation history [rom local observations and in doing so to gain further insight into how galaxies might have formed and evolved., The objective of this work is to reconstruct the cosmic star formation history from local observations and in doing so to gain further insight into how galaxies might have formed and evolved.494 As will be shown. this reconstruction is successful. comparing Tavorably with that deduced from high redshift observations (the Madau plot).," As will be shown, this reconstruction is successful, comparing favorably with that deduced from high redshift observations (the Madau plot)."495 This implies that our moclel provides a scenario for the chemical evolution of a representative sample of the universe., This implies that our model provides a scenario for the chemical evolution of a representative sample of the universe.496 Further. a proposed physical Braanework incorporates (he results from this chemical evolution model and predicts a major new stellar halo component. which we identilv in the Milkv Way with the (hick disk.," Further, a proposed physical framework incorporates the results from this chemical evolution model and predicts a major new stellar halo component, which we identify in the Milky Way with the thick disk."497 Classical elobular clusters like M92 belong to a group of relatively. blue. metal poor clusters which are old ancl are found in the outer halo of the Galaxy.," Classical globular clusters like M92 belong to a group of relatively blue, metal poor clusters which are old and are found in the outer halo of the Galaxy."498 It was early work by Becker (1950). Daade (1958). Morgan (1959) and Ninman (1959) which directed. attention to the disk-like spatial distribution of a second group of globular clusters which are reddcer. more metal rich. some what vounger. ancl concentrated towards the center of the Galaxy.," It was early work by Becker (1950), Baade (1958), Morgan (1959) and Kinman (1959) which directed attention to the disk-like spatial distribution of a second group of globular clusters which are redder, more metal rich, some what younger, and concentrated towards the center of the Galaxy."499 An important further difference is that whereas the metal poor group shows very little rotation about the Galactic center. the metal rich eroup shows significant rotation at all Galactocentrie radii (Cote. 1999. Zinn 1985).," An important further difference is that whereas the metal poor group shows very little rotation about the Galactic center, the metal rich group shows significant rotation at all Galactocentric radii (Cote, 1999, Zinn 1985)."500 A clear separation of the two groups shows in the frequency histogram of [Fe/IH] values for 133 Milky Way globular clusters in Fig., A clear separation of the two groups shows in the frequency histogram of [Fe/H] values for 133 Milky Way globular clusters in Fig.501 8 of Cote (1999) along wilh best fit Gaussians lor each group., 8 of Cote (1999) along with best fit Gaussians for each group.502 Recent spectroscopic studies of the globular clusters in M31 by Perrett et ((2002). and earlier work by Barmby et ((2000) and Huehra et ((1991) also show a bimodal distribution in Ρο with Gaussian parameters similar to those of the Milky Way.," Recent spectroscopic studies of the globular clusters in M31 by Perrett et (2002), and earlier work by Barmby et (2000) and Huchra et (1991) also show a bimodal distribution in [Fe/H] with Gaussian parameters similar to those of the Milky Way."503 In, In504of the luminosity [unction compared to when H(z)«1. indicating many more quasars just below the detection threshold (hat can be lensed into the sample.,"of the luminosity function compared to when $R(z) < 1$, indicating many more quasars just below the detection threshold that can be lensed into the sample."505 This leads (to an interesting conclusion: by setting the survey [lux limit artificially higher. one can expect to see a greater percentage of microlensed quasars.," This leads to an interesting conclusion: by setting the survey flux limit artificially higher, one can expect to see a greater percentage of microlensed quasars."506 A dramatic illustration of this is shown in Figure 5. for whieh the survey limit was made 1 mag brighter relative to Figure 4.., A dramatic illustration of this is shown in Figure \ref{fig:ew-dist-delta2} for which the survey limit was made 1 mag brighter relative to Figure \ref{fig:ew-dist-delta}.507 Everv 1 mag change in the survey limit multplies (2) by a [actor of 2.51: given the EDR. ancl 24Γ parameters from before. this means the new fIux ratio becomes larger than 1 for z>0.5 instead of lor 2>2.," Every 1 mag change in the survey limit multiplies $R(z)$ by a factor of 2.51; given the EDR and 2dF parameters from before, this means the new flux ratio becomes larger than 1 for $z > 0.5$ instead of for $z > 2$."508 The tradeolff. of course. is a reduced number of total observations as well as a smaller average optical depth to lensing for the sample.," The tradeoff, of course, is a reduced number of total observations as well as a smaller average optical depth to lensing for the sample."509 While the o-Dunction distribution is a useful beginning. a much more realistic model lor intrinsic equivalent widths is the lognormal distribution. llere (he “shape parameters” w and 5 are (he mean and standard deviation. respectively. of InWW.," While the $\delta$ -function distribution is a useful beginning, a much more realistic model for intrinsic equivalent widths is the lognormal distribution, Here the “shape parameters” $\omega$ and $\gamma$ are the mean and standard deviation, respectively, of $\ln W$."510 This is largely an empirically determined distribution. although there have been attempts to providea physical explanation (?)..," This is largely an empirically determined distribution, although there have been attempts to providea physical explanation \citep{qso:ew-model}."511 As will be shown in Section 3.. the lognormal distribution provides a very good approximation to the actual quasar data.," As will be shown in Section \ref{cha:data}, the lognormal distribution provides a very good approximation to the actual quasar data."512" Figure G shows how the lensed equivalent width model appears for various values of z and ο), and choosing w and 5 as appropriate for the data set."," Figure \ref{fig:ew-dist-lognormal} shows how the lensed equivalent width model appears for various values of $z$ and $\Omega_c$, and choosing $\omega$ and $\gamma$ as appropriate for the data set."513 As in Figure 4.. there is a clear lensing indicator in the low equivalent width region that scales with both z and Ως.," As in Figure \ref{fig:ew-dist-delta}, , there is a clear lensing indicator in the low equivalent width region that scales with both $z$ and $\Omega_c$."514 Because of the laree amount of integration involved in arriving at (he final expression (28)) for the lensed-lognormal equivalent width distribution. it becomes necessary (to compute an interpolated version of the model.," Because of the large amount of integration involved in arriving at the final expression \ref{eq:pW}) ) for the lensed-lognormal equivalent width distribution, it becomes necessary to compute an interpolated version of the model."515 Through a change of variables. 1 is possible to remove w dependence [rom (he integral of the lensed-lognormal model.," Through a change of variables, it is possible to remove $\omega$ dependence from the integral of the lensed-lognormal model."516 Then polynomial interpolation of the expression over the remaining four variables (4.2.5.£2.) gives a very good fit to the actual distribution.," Then polynomial interpolation of the expression over the remaining four variables $(u, z, \gamma, \Omega_c)$ gives a very good fit to the actual distribution."517 In the £9).=(0 case. Equation 31.is a simple parabola in 4. independent of all other parameters.," In the $\Omega_c = 0$ case, Equation \ref{eq:pW-xform} is a simple parabola in $u$ , independent of all other parameters."518of our model inadequacy criterion in the Appendix.,of our model inadequacy criterion in the Appendix.519 Finally. in section 3 we apply this criterion in practice by analysing astronomical RV exoplanet detections made using at least two different telescope-instrument combinations.," Finally, in section 3 we apply this criterion in practice by analysing astronomical RV exoplanet detections made using at least two different telescope-instrument combinations."520 The Bayesian methods do not differentiate between determining the most probable parameter values or most probable models containing these parameters., The Bayesian methods do not differentiate between determining the most probable parameter values or most probable models containing these parameters.521 They can all be arranged into a linear order. which yields information on the observed system if only the selected models describe the observed system realistically enough.," They can all be arranged into a linear order, which yields information on the observed system if only the selected models describe the observed system realistically enough."522 It is possible to calculate the relative posterior probabilities of any number of models and determine their relative magnitudes mn a similar way as it is possible to determine the posterior odds of having the measurements drawn from a probability density characterised by a certain parameter value of any one of the models., It is possible to calculate the relative posterior probabilities of any number of models and determine their relative magnitudes in a similar way as it is possible to determine the posterior odds of having the measurements drawn from a probability density characterised by a certain parameter value of any one of the models.523 We do not deseribe the process of determining the posterior probability densities of the model parameters here. because several well-known posterior sampling methods exist and they have been well covered by the existing literature (e.g.1984:Haarioetal.. 2001).," We do not describe the process of determining the posterior probability densities of the model parameters here, because several well-known posterior sampling methods exist and they have been well covered by the existing literature \citep[e.g.][]{metropolis1953,hastings1970,geman1984,haario2001}."524. The performance of these methods has also been demonstrated by several re-analyses of existing RV data. revealing the existence of planets (e.g.Gregory. or disputing it (e.g.Tuomi. 2011).," The performance of these methods has also been demonstrated by several re-analyses of existing RV data, revealing the existence of planets \citep[e.g.][]{gregory2005,gregory2007a,gregory2007b,tuomi2009} or disputing it \citep[e.g.][]{tuomi2011}."525. In these works. the model probabilities have played an important role in assessing the number of planetary companions orbiting nearby stars.," In these works, the model probabilities have played an important role in assessing the number of planetary companions orbiting nearby stars."526 Commonly. the Bayesian tools are used to assess the probabilities of different statistical models given. the measurements 7; that are being analysed using the models.," Commonly, the Bayesian tools are used to assess the probabilities of different statistical models given the measurements $m$ that are being analysed using the models."527 These tools provide the relative probabilities of the selected models Afj;j=1....k in the determined model set as where probabilities P(At).i=]l....& are the prior probabilities of the different models and the marginal likelihoods POan|At;) are defined as where z(6;) is the prior probability density of the parameter or parameter vector 6; of the model At; and /65]|6;) represents the likelihood function corresponding to the model.," These tools provide the relative probabilities of the selected models $\mathcal{M}_{i}, i=1, ..., k$ in the determined model set as where probabilities $P(\mathcal{M}_{i}), i=1, ..., k$ are the prior probabilities of the different models and the marginal likelihoods $P(m | \mathcal{M}_{i})$ are defined as where $\pi(\theta_{i})$ is the prior probability density of the parameter or parameter vector $\theta_{i}$ of the model $\mathcal{M}_{i}$ and $l(m | \theta_{i})$ represents the likelihood function corresponding to the model."528 The interpretation of the posterior probabilities in Eq. (1)), The interpretation of the posterior probabilities in Eq. \ref{model_probability}) )529 is à rather subjective matter because they are relative and it is only possible to assess how much confidence one has in one of the models compared to the rest of them., is a rather subjective matter because they are relative and it is only possible to assess how much confidence one has in one of the models compared to the rest of them.530 According to the views of Jeffreys(1961):Kass&Raftery(1995).. a model would have to be at least 150 times more probable than the next best model to have strong evidence in favour of it.," According to the views of \citet{jeffreys1961,kass1995}, a model would have to be at least 150 times more probable than the next best model to have strong evidence in favour of it."531 We adopt the same threshold because claiming that there are &4| planets orbiting a star instead of k needs to be on a solid ground with respect to the model probabilities., We adopt the same threshold because claiming that there are $k+1$ planets orbiting a star instead of $k$ needs to be on a solid ground with respect to the model probabilities.532 Especially. if the model with k+1 planets was e.g. 50 times more probable than that with & planets. there would still be a roughly possibility that the & planet model explains the data.," Especially, if the model with $k+1$ planets was e.g. 50 times more probable than that with $k$ planets, there would still be a roughly possibility that the $k$ planet model explains the data."533 Therefore. we choose a rather high threshold when interpreting the posterior probabilities of models with different numbers of Keplerian signals.," Therefore, we choose a rather high threshold when interpreting the posterior probabilities of models with different numbers of Keplerian signals."534 With the marginal likelihoods available according to the Eq. 2..," With the marginal likelihoods available according to the Eq. \ref{marginal_likelihood},"535" we define the model Af to be an inadequate description of independent measurements 715.7=1...N. If it holds that for some small positive number + This definition isIh, based on the independence of the measurements and that they are being modelled with a single statistical model."," we define the model $\mathcal{M}$ to be an inadequate description of independent measurements $m_{i}, i=1, ..., N$, if it holds that for some small positive number $r$ This definition is based on the independence of the measurements and that they are being modelled with a single statistical model."536 It is a simple result of a relation of the marginal likelthoods of each of the measurement and the joint marginal likelihood of all of them shown in Eq. (A8))., It is a simple result of a relation of the marginal likelihoods of each of the measurement and the joint marginal likelihood of all of them shown in Eq. \ref{multiple_probabilities}) ).537 We derive this criterion using the Bayes’ rule of conditional probabilities and the concept of independence. and also interpret the results in terms of information theory in. the Appendix.," We derive this criterion using the Bayes' rule of conditional probabilities and the concept of independence, and also interpret the results in terms of information theory in the Appendix."538 The number + has an interpretation as a threshold value., The number $r$ has an interpretation as a threshold value.539 For instance. the model being inadequate with probabilities90%...955c.. and corresponds to threshold values of 0.111. 0.053. and 0.010. respectively (see Appendix).," For instance, the model being inadequate with probabilities, and corresponds to threshold values of 0.111, 0.053, and 0.010, respectively (see Appendix)."540 Therefore. if the best model according to Eq. (1))," Therefore, if the best model according to Eq. \ref{model_probability}) )"541 satisfies Eq. (3)), satisfies Eq. \ref{MIC}) )542 for some reasonably small +. it can be concluded that the model does not describe the measurements without bias and the corresponding analysis results may be biased as well.," for some reasonably small $r$, it can be concluded that the model does not describe the measurements without bias and the corresponding analysis results may be biased as well."543 In such a case. the model set has to be re-considered and expanded by adding better descriptions of the data to it.," In such a case, the model set has to be re-considered and expanded by adding better descriptions of the data to it."544 In practice. we use the threshold value. but choosing its value is a subjective issue and only represents how confidently one wants to determine the model inadequacy.," In practice, we use the threshold value, but choosing its value is a subjective issue and only represents how confidently one wants to determine the model inadequacy."545 We note that the model inadequacy can also be interpreted in terms of the measurements being inconsistent with one another with respect to the model used., We note that the model inadequacy can also be interpreted in terms of the measurements being inconsistent with one another with respect to the model used.546 This interpretation arises from the fact that the model may not take into account some features in one or more data sets that result from biases in the process of making the measurements or from some other unmodelled features in the data., This interpretation arises from the fact that the model may not take into account some features in one or more data sets that result from biases in the process of making the measurements or from some other unmodelled features in the data.547 We use the inadequacy of the model given the data sets and the inconsistency of the data sets with respect to this model interchangeably throughout this article., We use the inadequacy of the model given the data sets and the inconsistency of the data sets with respect to this model interchangeably throughout this article.548 We describe the parameter probability densities using three numbers., We describe the parameter probability densities using three numbers.549 These numbers are the maximum (MAP) estimate of the posterior density and the limits of the Bayesian credibility set 5055s as defined in e.g. Tuomi&Koti-ranta (2009)., These numbers are the maximum (MAP) estimate of the posterior density and the limits of the Bayesian credibility set $\mathcal{D}_{0.99}$ as defined in e.g. \citet{tuomi2009}.550. We calculate these estimates from the posterior densities of the model parameters received using the adaptive Metropolis posterior sampling algorithm (Haarioetal..2001).. which is a modification of the famous Metropolis-Hastings (M-H) algorithm (Metropolisetal..1953:Hastings.1970). that adapts the proposal density to the shape of the posterior density of the model parameters.," We calculate these estimates from the posterior densities of the model parameters received using the adaptive Metropolis posterior sampling algorithm \citep{haario2001}, which is a modification of the famous Metropolis-Hastings (M-H) algorithm \citep{metropolis1953,hastings1970} that adapts the proposal density to the shape of the posterior density of the model parameters."551 Because of this property. it is not very sensitive to the choise of initial parameter vector nor proposal density — desired features that make the method significantly," Because of this property, it is not very sensitive to the choise of initial parameter vector nor proposal density – desired features that make the method significantly"552 , 553when fitting the nuclear integrated spectrum of this source within a larger aperture (r~200 ppc) over the sspectral region.,when fitting the nuclear integrated spectrum of this source within a larger aperture $\simeq$ pc) over the spectral region.554 We point out that the synthetic spectra was reddened using the Ay derived from the synthesis (Fig. 7))., We point out that the synthetic spectra was reddened using the $\rm_V$ derived from the synthesis (Fig. \ref{pop2}) ).555" Following Riffeletal.(2010b,seealsoal.,2009andCid-Fernandesetal. 2004),, we have binned the contribution of the SPCs (xj) into a reduced population vector with four age ranges:young (x,: t<100 Myr); (x: 0.3<t« Gyr), (xi; 1«t2 Gyr) and (x: 5«t13 Gyr)."," Following \citet[][see also Riffel et al., 2009 and Cid-Fernandes et556al. 2004]{rogemar10b}, we have binned the contribution of the SPCs $x_j$ ) into a reduced population vector with four age ranges: $x_y$ : $t \leq 100$ Myr); $x_{yi}$: $0.3 \leq t \leq5570.7$ Gyr), $x_{io}$: $1 \leq t \leq 2$ Gyr) and $x_{o}$: $5 \leq t \leq 13$ Gyr)."558" In Fig. 5,,"," In Fig. \ref{pop1},"559" we show the spatial distribution of the percent flux contribution at 2.12 of the stars in each X: while in regions farther than r~ umppc) from the nucleus the contribution of the young-intermediate SPCs reaches values of up to100%,, closer to the nucleus the contribution of this component is negligible."," we show the spatial distribution of the percent flux contribution at $\mu$ m of the stars in each $\vec{x}$: while in regions farther than $\sim$ pc) from the nucleus the contribution of the young-intermediate SPCs reaches values of up to, closer to the nucleus the contribution of this component is negligible."560" Within this region the stellar population is dominated by the old component(~50%)), closely surrounded (r x 0744) by an intermediate-old population (also ~50%))."," Within this region the stellar population is dominated by the old $\approx$ ), closely surrounded (r $\lesssim$ 4) by an intermediate-old population (also $\approx$ )."561" No signs of young stellar populations were detected close to the nucleus; however, a significant contribution (25-60%)) is found in regions farther than r 0.66."," No signs of young stellar populations were detected close to the nucleus; however, a significant contribution ) is found in regions farther than $r\sim$ 6."562 The light-fraction SPC contributions depend of the normalization wavelength and thus the comparison with results from other spectral regions should be done with caution (Riffeletal. 2010c)., The light-fraction SPC contributions depend of the normalization wavelength and thus the comparison with results from other spectral regions should be done with caution \citep{rogerio10}.563". However, a physical parameter which does not depend on the normalization point used in the fit is the stellar mass."," However, a physical parameter which does not depend on the normalization point used in the fit is the stellar mass."564 The mass-fraction of each population vector components is show in Fig., The mass-fraction of each population vector components is show in Fig.565" 5 (young: young-intermediate: m,;, intermediate-old: mio and old population:m,, m,)."," \ref{pop1} (young: $m_y$, young-intermediate: $m_{yi}$, intermediate-old: $m_{io}$ and old population: $m_o$ )."566 The maps of the mass-weighted SPC follow a similar distribution to the light-weighted ones., The maps of the mass-weighted SPC follow a similar distribution to the light-weighted ones.567" However, in the former the contribution of the older ages is enhanced, particularly within 0744 from the nucleus."," However, in the former the contribution of the older ages is enhanced, particularly within 4 from the nucleus."568" Besides the SPC distributions, the ooutputs the average reddening of the stellar populations refpop2))."," Besides the SPC distributions, the outputs the average reddening of the stellar populations )."569 The highest values of E(B—V)=0.7 (we used Ay=3.1E(B-V)) are reached at the nucleus up to within «0:22 from it., The highest values of $E(B-V)=0.7$ (we used $\rm A_V$ =3.1E(B-V)) are reached at the nucleus up to within $\sim$ 2 from it.570" The goodness of the fit is measured in bby the percent mean deviation: adev==|O,—M;|/O;, where Οι is the observed spectrum and M, is the fitted model (CidFernandesetal.2004,2005b)."," The goodness of the fit is measured in by the percent mean deviation: $|O_\lambda-M_\lambda|/O_\lambda$, where $O_\lambda$ is the observed spectrum and $M_\lambda$ is the fitted model \citep{cid04,cid05}."571. The map for 11157 is shown in Fig., The map for 1157 is shown in Fig.572 7 and presents values adevx5 well the observed underlying spectra., \ref{pop2} and presents values $adev\lesssim5$ well the observed underlying spectra.573" We point out that, in order to have a robust result on stellar population fitting with iit is important to have a reliable flux calibration (see mmanual)."," We point out that, in order to have a robust result on stellar population fitting with it is important to have a reliable flux calibration (see manual)."574" In other words the fit depends on the overall shape of the observed spectrum from the J to the K-band, and our NIFS observations misses the H-band."," In other words the fit depends on the overall shape of the observed spectrum from the $J$ to the $K$ -band, and our NIFS observations misses the $H$ -band."575" In order to verify the reliability of our flux calibration we have extracted a spectrum from our datacubes matching the aperture and position angle of our SpeX cross-dispersed spectrum (Riffel,Rodríguez-Ardila,2006a).", In order to verify the reliability of our flux calibration we have extracted a spectrum from our datacubes matching the aperture and position angle of our SpeX cross-dispersed spectrum \citep{rogerio06}.576". Note that the SpeX data were taken in the cross-dispersed mode, and thus are free from seeing and aperture effects (seeRiffel,Rodríguez-Ardila,&Pastoriza2006a,for details).."," Note that the SpeX data were taken in the cross-dispersed mode, and thus are free from seeing and aperture effects \citep[see][for details]{rogerio06}."577 The comparison between both spectra is shown in Fig. 4.., The comparison between both spectra is shown in Fig. \ref{nifsxspex}.578 It is clear that our NIFS spectra havea reliable flux calibration., It is clear that our NIFS spectra have a reliable flux calibration.579 In fact the difference between the NIFS and SpeX spectra is lower than in the J band., In fact the difference between the NIFS and SpeX spectra is lower than in the $J$ band.580" In addiditon, we have performed a simulation varying the flux of the J band by (from up to 110%)) and the difference to the population vector components is, in mean, $5% indicating that our results are robust."," In addiditon, we have performed a simulation varying the flux of the $J$ band by (from up to ) and the difference to the population vector components is, in mean, $\lesssim$ indicating that our results are robust."581" In general, our results are similar to those obtained in the 2D mapping of the SP of 11066 (Riffeletal.2010b) as well as for previous NIR studies using single aperture nuclear spectra (Riffeletal. 2009d).."," In general, our results are similar to those obtained in the 2D mapping of the SP of 1066 \citep{rogemar10b} as well as for previous NIR studies using single aperture nuclear spectra \citep{rogerio09}. ."582 By mapping the stellar population in 2D we can analyse the spatial variations of the SPCs in the inner few hundred parsecs of the galaxy., By mapping the stellar population in 2D we can analyse the spatial variations of the SPCs in the inner few hundred parsecs of the galaxy.583" Further, the significance of these variations is enhanced by the comparison with the σ. map presented in Fig. 6.."," Further, the significance of these variations is enhanced by the comparison with the $\sigma_*$ map presented in Fig. \ref{sig}."584" This map was obtained by fitting the K-band CO absorption heads with the penalized Pixel Fitting (pPXF) method of Cappellari&Emsellem(2004) using as stellar template spectra those of the Gemini library of late spectral type stars observed with the Gemini Near-Infrared Spectrograph (GNIRS) IFU and NIFS (Winge,Riffel&Storchi-Bergmann 2009).", This map was obtained by fitting the K-band CO absorption band-heads with the penalized Pixel Fitting (pPXF) method of \citet{cappellari04} using as stellar template spectra those of the Gemini library of late spectral type stars observed with the Gemini Near-Infrared Spectrograph (GNIRS) IFU and NIFS \citep{winge09}.585. More details on the stellar kinematics of the central region of 11157 can be found in Riffel prep.)., More details on the stellar kinematics of the central region of 1157 can be found in Riffel ).586" The c"". map shows partial ring of low-c-, values (5 kms!) surrounding the anucleus at ~ 0/88 ppc) from it, immersed in higher c, values of the bulge stars (« km s!)."," The $\sigma_*$ map shows a partial ring of $\sigma_*$ values $\approx$ ${\rm km\,s^{-1}}$ ) surrounding the nucleus at $\approx$ 8 pc) from it, immersed in higher $\sigma_*$ values of the bulge stars $\approx$ ${\rm km\,s^{-1}}$ )."587" Such rings are commonly observed in the central region of active galaxies and are due to kinematically colder regions with younger stars than the underlyingbulge(Barbosaetal.2006;Deo, 2009a).."," Such rings are commonly observed in the central region of active galaxies and are due to kinematically colder regions with younger stars than the underlyingbulge\citep{barbosa06,deo06,lopes07,rogemar08,rogemar09a}. ."588" The comparison between 11157 stellar population synthesis maps (light- and mass-weighted) and the c. map shows that the low o. ring is spatially correlated with the young-intermediate age SPC, while the highest os are associated with the old component."," The comparison between 1157 stellar population synthesis maps (light- and mass-weighted) and the $\sigma_*$ map shows that the low $\sigma_*$ ring is spatially correlated with the young-intermediate age SPC, while the highest $\sigma_*$ s are associated with the old component."589" Interestingly, very similarresults were found by Riffeletal.(2010b) for 11066."," Interestingly, very similarresults were found by \citet{rogemar10b} for 1066."590" Thus, the results found for 11157 and 11066 support the use of low stellar"," Thus, the results found for 1157 and 1066 support the use of low stellar"591as described in LOs.,as described in L08.592 Tere we extend this work aud analyze independeut dynamical information on the ealaxy velocities aud also the number deusity profile of cluster iienibers., Here we extend this work and analyze independent dynamical information on the galaxy velocities and also the number density profile of cluster members.593 Iu this section we explain in detail our analysis of the latter two data sets., In this section we explain in detail our analysis of the latter two data sets.594 Establishing the form of the projected profile of cluster ucniber galaxies is essential for fully cimplovine the Jeans equation (Binney Tremaine 1987: see 3)). which allows for a spatial distribution of galaxies that 1ced not follow the dominant DM.," Establishing the form of the projected profile of cluster member galaxies is essential for fully employing the Jeans equation (Binney Tremaine 1987; see \ref{Methodology}) ), which allows for a spatial distribution of galaxies that need not follow the dominant DM."595 \leasurcement of the sojected galaxy distribution requires subtraction of the vackeround aud foreground field galaxy populations., Measurement of the projected galaxy distribution requires subtraction of the background and foreground field galaxy populations.596 This must be achieved with accurate multi-color photometry as spectroscopy is usually Πίος to πια. siuuples aud docs not extend faint enough to include the uajoritv of cluster menibers., This must be achieved with accurate multi-color photometry as spectroscopy is usually limited to small samples and does not extend faint enough to include the majority of cluster members.597 We used Subaru photometry in the V aud I bauds., We used Subaru photometry in the V and I bands.598" The data are colplete to a depth of Ly,=26.5. reaching cight magnitudes below L of the luminosity function."," The data are complete to a depth of $I_{AB}=26.5$, reaching eight magnitudes below $L^{*}$ of the luminosity function."599 We selected a color region based on a color-1nagnitude relation analysis. which comprises 17171 galaxies.," We selected a color region based on a color-magnitude relation analysis, which comprises 17474 galaxies."600 This color region ranges from the red side of the E/SO cluster sequence to a blue boundary chosen so that the sample exteuds sufficiently. blueward to include most (~ 85%) of the cluster members as described in \ledezinski ct ((2007. hereafter MOT. fie.," This color region ranges from the red side of the E/SO cluster sequence to a blue boundary chosen so that the sample extends sufficiently blueward to include most $\sim 85\%$ ) of the cluster members as described in Medezinski et (2007, hereafter M07, fig."601 1)., 1).602 MOT established that this color region contains the majority of cluster nienibers iu addition to some background galaxies. by exaiuiniue the weak lensing signal.," M07 established that this color region contains the majority of cluster members in addition to some background galaxies, by examining the weak lensing signal."603 In this color region. which iucludes the E/SO sequence and bluer objects Guclidiug some backgrouud galaxies) the lensing signal was found to be sienificautly lower than the true background signal as measured for red backeround galaxies.," In this color region, which includes the E/SO sequence and bluer objects (including some background galaxies) the lensing signal was found to be significantly lower than the true background signal as measured for red background galaxies."604 This is because (uuleused) cluster 1ieuibers dilute the weal leusiug signal of the background. allowing us to ideutifv the region of color space occupied by cluster ΠΟΙΟΥΣ.," This is because (unlensed) cluster members dilute the weak lensing signal of the background, allowing us to identify the region of color space occupied by cluster members."605 We derived the projected galaxy distribution of cluster ucnibers from the above colorselected galaxy sample vased on fitting the radial distribution., We derived the projected galaxy distribution of cluster members from the above color-selected galaxy sample based on fitting the radial distribution.606 Calaxies were first raclially binned iuto anu aud the surface nuuber density in each annulus was determined: uote that at aree radi onlv part of each annulus was covered by he detector., Galaxies were first radially binned into annuli and the surface number density in each annulus was determined; note that at large radii only part of each annulus was covered by the detector.607 Assuming Poissonian (IN) crrors. the ealaxv surface number density was modelled as that due to the cluster plus a background galaxy surface προς deusitv. with the latter assumed muiform.," Assuming Poissonian $\sqrt{N}$ ) errors, the galaxy surface number density was modelled as that due to the cluster plus a background galaxy surface number density, with the latter assumed uniform."608 Thus. we effectively. neelected. fluctuations in the backerouud ealaxy surface nuniber deusity on scales simaller the cluster size (about 22/ at 2= 0.183).," Thus, we effectively neglected fluctuations in the background galaxy surface number density on scales smaller the cluster size (about $22^\prime$ at $z=0.183$ )."609 There could be such fluctuations at some level due to correlated structures (such as filaments) along the line of sight to the cluster., There could be such fluctuations at some level due to correlated structures (such as filaments) along the line of sight to the cluster.610 Calaxies in these structures could mistakenly be imcluded as cluster ΠΟΙΟΥΣ. a possibility that we assess below.," Galaxies in these structures could mistakenly be included as cluster members, a possibility that we assess below."611 Figure shows that the radial profile of the above color-sclected galaxy sample is well fitted with a wniform backerounud plus a general cored profile to represent the cluster galaxy surface density profile: where M4 ix the total ealaxy surface uuuber deusity. Xo. re and p are the three parameters of the cored profile. ad C5. is the background density.," Figure \ref{galaxy surface number density} shows that the radial profile of the above color-selected galaxy sample is well fitted with a uniform background plus a general cored profile to represent the cluster galaxy surface density profile: where $\Sigma_{\rm tot}$ is the total galaxy surface number density, $\Sigma_0$, $r_c$ , and $p$ are the three parameters of the cored profile, and $C_{\rm bg}$ is the background density."612 The resulting fit is eood. \?/dof=19.2/16.," The resulting fit is good, $\chi^2/{\rm dof}=19.2/16$ ."613 Using a larger umber of bins than No=20 improves the reduced 47. ea. NM=50 aud N=100 eive \?/dof=36/16 aud \?/dof=96/96. respectively.," Using a larger number of bins than $N=20$ improves the reduced $\chi^2$, e.g., $N=50$ and $N=100$ give $\chi^2/{\rm dof}=36/46$ and $\chi^2/{\rm dof}=96/96$ , respectively."614 The level of the background is stable aud does not depend on the nuuber of bius. e.g.. for Vo20. 50. aud 100 we find C=LOTL+10. LOGL+12. and 1063=LL. respectively.," The level of the background is stable and does not depend on the number of bins, e.g., for $N=20$ , 50, and 100 we find $C=1071\pm 40$, $1064 \pm 42$, and $1063\pm 41$, respectively."615 We settled ou just 20 bius. since this made it easier to combine the surface deusity data set with the projected velocity dispersion data set. which is muferior m terms of signal to noise per radial bin (see below).," We settled on just 20 bins, since this made it easier to combine the surface density data set with the projected velocity dispersion data set, which is inferior in terms of signal to noise per radial bin (see below)."616 With 20 bius. we fom best-fit values of My1200+110 b? 727 =0+150 3 kpe. and p=0.7LEzE0.30.," With 20 bins, we found best-fit values of $\Sigma_0=1200 \pm 140$ $^2$ $^{-2}$, $r_c = 450 \pm 150$ $^{-1}$ kpc, and $p=0.74 \pm 0.30$ ."617 To ineasure the iass profile of galaxy clusters using ealaxy motions it is necessary to obtain precise velocity iieasurenments for a statistically large siuuple of ealaxies., To measure the mass profile of galaxy clusters using galaxy motions it is necessary to obtain precise velocity measurements for a statistically large sample of galaxies.618 The data used here are part of an extensive iulti-object spectroscopy survey carried out with the VIMOS spectrograph ou the VET. (Czoske 2001): for observational details. see Czoske (2001).," The data used here are part of an extensive multi-object spectroscopy survey carried out with the VIMOS spectrograph on the VLT (Czoske 2004); for observational details, see Czoske (2004)."619 This dataset constitutes 1169. objects with reliable spectroscopic redshifts. a quajor advance over previous surveys of Al6s9.," This dataset constitutes 1469 objects with reliable spectroscopic redshifts, a major advance over previous surveys of A1689."620 Note though that we did not try to find the ealaxy surface density profile from the projected velocity data set. since the data set used in the previous section contains a much larecr umber of galaxies.," Note though that we did not try to find the galaxy surface density profile from the projected velocity data set, since the data set used in the previous section contains a much larger number of galaxies."621 We analyzed the spectroscopic sample by first defining cluster membership using the velocity “caustics”. which are clearly visible for this cluster iu the form of a boundary which varies with radius. peaking at arouud +1000 ins at ~300 hi. kpce aud declining steadily at larger radius (see figure 2)).," We analyzed the spectroscopic sample by first defining cluster membership using the velocity “caustics”, which are clearly visible for this cluster in the form of a boundary which varies with radius, peaking at around $\pm 4000$ km/s at $\sim 300$ $^{-1}$ kpc and declining steadily at larger radius (see figure \ref{velocity space622diagram}) )."623 The caustics are related to the escape velocity from the clusteraud thus provide a tangibleplivsical basis by which we cau separate cluster members from foreground. and background galaxies., The caustics are related to the escape velocity from the clusterand thus provide a tangiblephysical basis by which we can separate cluster members from foreground and background galaxies.624 Defining membershipis especially important for. massive, Defining membershipis especially important for massive625bv the DOE-supported ASC/Alliauce.— Ceuter for Astrophysical Thermonuclear| Flashes at the University of Chicago.,by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.626 MIR thauks Jeremy Wallin for lis invaluable help with maintaining the computing cluster at the Michigan Academic Compuine Ceuter where most of the computatious were performed., MR thanks Jeremy Hallum for his invaluable help with maintaining the computing cluster at the Michigan Academic Computing Center where most of the computations were performed.627 MB thanks Justin Nieusmia for technical assisance with performing the smunlatious., MR thanks Justin Nieusma for technical assistance with performing the simulations.628 We thauk Dorewook Lee. Iun Parrish. Eliot Quatacrt. Elena Rasia. Prateck Sharia and Miu-Su Shin for discussious.," We thank Dongwook Lee, Ian Parrish, Eliot Quataert, Elena Rasia, Prateek Sharma and Min-Su Shin for discussions."629 We would like to point out that couclusious from cooiip runs. simular tfo the ones preseuted here. were «jbtained independently by Iu Parrish and collaboraors.," We would like to point out that conclusions from cooling runs, similar to the ones presented here, were obtained independently by Ian Parrish and collaborators."630 We iue erateful to this eroup for sluuiug some of thei results., We are grateful to this group for sharing some of their results.631 SPO acknowledges support by NASA erant. NNCGOGCII9SC:. and NSF οτα 0908[80.," SPO acknowledges support by NASA grant NNG06GH95G, and NSF grant 0908480."632. AIR acknowledges support by theory eraut. TM8S-9011X. MB and SPO thaux Iustitute of Astronomy. Cambridge. UR and Max Pluck Iustitute for Astroplivsies. Garching. Germany for their hospitality.," MR acknowledges support by theory grant TM8-9011X. MR and SPO thank Institute of Astronomy, Cambridge, UK and Max Planck Institute for Astrophysics, Garching, Germany for their hospitality."633Since their discovery thirty vears ago 11963). the nature and origin of high velocity clouds (1VCSs) has remained highly controversial.,"Since their discovery thirty years ago 1963), the nature and origin of high velocity clouds (HVCs) has remained highly controversial."634 Vheir checquerecl history has been discussed by Verschuur (1988) and Wakker van Woerden (1997)., Their chequered history has been discussed by Verschuur (1988) and Wakker van Woerden (1997).635 LIVCs which cover at least a third of the sky are concentrations of neutral hydrogen with velocities which do not conform to a simple mocel of galactic rotation., HVCs $-$ which cover at least a third of the sky $-$ are concentrations of neutral hydrogen with velocities which do not conform to a simple model of galactic rotation.636 Few. io any. clouds have reliable clistance cleterminations. which has encouraged. wide ranging speculation as to their origin.," Few, if any, clouds have reliable distance determinations, which has encouraged wide ranging speculation as to their origin."637 Explanations (assumed clistances are given in parentheses) range from. local supernova remnants. (7100 20). large-scaleὃν expandinge motions in nearby spiral armis (« 1 kpe). condensations in the local galactic halo (~ 1 kpe). structures in the ealactic warp 20 kpe). tidal clisruptions of the Magellanie Clouds 50 kpe). intergalactic gas (2 50 kpe) or protogalaxies (~500 kpc).," Explanations (assumed distances are given in parentheses) range from local supernova remnants $\sim$ 100 pc), large-scale expanding motions in nearby spiral arms $<$ 1 kpc), condensations in the local galactic halo $\sim$ 1 kpc), structures in the galactic warp $-$ 20 kpc), tidal disruptions of the Magellanic Clouds $-$ 50 kpc), intergalactic gas $>$ 50 kpc) or protogalaxies $\sim$ 500 kpc)."638 The distance uncertainty continues to be the major stumbling block in understanding HIVC€s (Schwarz. Wakker van Woerden 1995) since the cloud density and mass scale inversely. with distance. and as the square of the distance. respectively.," The distance uncertainty continues to be the major stumbling block in understanding HVCs (Schwarz, Wakker van Woerden 1995) since the cloud density and mass scale inversely with distance, and as the square of the distance, respectively."639 Ferrara Field (1994: see also 11995) have suggested one possible method for svstematic distance. determinations based. on. μα core/envelope structure observed in some LIVCs (Cram Ciovanelli 1976)., Ferrara Field (1994; see also 1995) have suggested one possible method for systematic distance determinations based on the core/envelope structure observed in some HVCs (Cram Giovanelli 1976).640 Part of the problem is that. besides oobservations. it has proved difficult to detect. ΕΝΟΣ in," Part of the problem is that, besides observations, it has proved difficult to detect HVCs in"641formation. the dust grain growth: would always proceed in a hot environment.,"formation, the dust grain growth would always proceed in a hot environment."642 Then the resulting planetesimals would be devoid of water. (, Then the resulting planetesimals would be devoid of water. (6432) Uniform dust grain size and cdust-to-gas mass ratio: In a real disk. the dust grain size and the dust-to-gas mass ratio would be non-uniform.,"2) Uniform dust grain size and dust-to-gas mass ratio: In a real disk, the dust grain size and the dust-to-gas mass ratio would be non-uniform."644 Dust particles grow mainly bv collisions., Dust particles grow mainly by collisions.645 Ànd the growth tends to proceed [rom the inner disk region. so the dust grain size in the inner disk is likelv to be larger than that in the outer disk.," And the growth tends to proceed from the inner disk region, so the dust grain size in the inner disk is likely to be larger than that in the outer disk."646 When the cust grain is large enough. the height of the disk surface is lowered.," When the dust grain is large enough, the height of the disk surface is lowered."647 Consequently. a disk region. which is located. just outside the developed dust erain region aud the dust grain growth: is about lo start. can receive more stellar radiation flux.," Consequently, a disk region, which is located just outside the developed dust grain region and the dust grain growth is about to start, can receive more stellar radiation flux."648 Then. the dust grain growth may proceed with a hot temperature aud water deficit planetesimals may form. (," Then, the dust grain growth may proceed with a hot temperature and water deficit planetesimals may form. ("6493) Solid material and water distribution evolution: According to Ciesla Cuzzi (2006) and Garaud (2007). à large amount of water vapor and fine dust particles are (ransported from the outer part to the inner part of the disk during the disk evolution.,"3) Solid material and water distribution evolution: According to Ciesla Cuzzi (2006) and Garaud (2007), a large amount of water vapor and fine dust particles are transported from the outer part to the inner part of the disk during the disk evolution."650 Η fine dust parücles or water vapor are supplied to the terrestrial planet. region. ancl the snow line location is kept farther from (he terrestrial planet region until the planetesimal formation is completed. the dust grain. growth may proceed in a water ice-free environment. and consequently water-devoid planetesimals may be formed.," If fine dust particles or water vapor are supplied to the terrestrial planet region, and the snow line location is kept farther from the terrestrial planet region until the planetesimal formation is completed, the dust grain growth may proceed in a water ice-free environment, and consequently water-devoid planetesimals may be formed."651 Whether or not this mechanism works may depend on the initial condition of the disk., Whether or not this mechanism works may depend on the initial condition of the disk.652 Yel another possible scenario may be an elimination of water from icv. planetesinals alter the snow line leaves outwardly (he terrestrial planet region., Yet another possible scenario may be an elimination of water from icy planetesimals after the snow line leaves outwardly the terrestrial planet region.653 Future work on these points is needed., Future work on these points is needed.654 In this study. it is assumed (hat ice condenses as pure ice particles and no ice mantle," In this study, it is assumed that ice condenses as pure ice particles and no ice mantle"655Galaxy clusters have long been thought of as particularly active sites of galaxy evolution (2)..,Galaxy clusters have long been thought of as particularly active sites of galaxy evolution \citep{Dressler1980}.656 Early optical studies suggested that clusters affect stellar populations in infalling substructure mainly by turning off star formation as galaxies encounter the dense cluster environment (?).., Early optical studies suggested that clusters affect stellar populations in infalling substructure mainly by turning off star formation as galaxies encounter the dense cluster environment \citep{Lewis2002}.657 However. when these same objects were observed in the infrared by IRAS. ISO andSpitzer.. a significant fraction of cluster galaxies was revealed to contain heavily obscured sites of star formation (see e.g. ? and references therein).," However, when these same objects were observed in the infrared by IRAS, ISO and, a significant fraction of cluster galaxies was revealed to contain heavily obscured sites of star formation (see e.g. \citet{Metcalfe2005} and references therein)."658 When this obscured activity 1s accounted for. clusters emerge as environments where star formation is both quenched and induced. perhaps even at different stages of the same physical process.," When this obscured activity is accounted for, clusters emerge as environments where star formation is both quenched and induced, perhaps even at different stages of the same physical process."659 Some studies have claimed to observe triggered star formation in strongly disrupted. merging clusters (e.g.?) but the intrinsic scatter in this correlation is large. with some seemingly relaxed clusters presenting much higher star formation activity than their actively merging counterparts (?)..," Some studies have claimed to observe triggered star formation in strongly disrupted, merging clusters \citep[e.g.][]{Geach2006} but the intrinsic scatter in this correlation is large, with some seemingly relaxed clusters presenting much higher star formation activity than their actively merging counterparts \citep{Haines2009}."660 Many studies in recent years have attempted to separate cluster samples into two distinct populations: merging clusters vs clusters that are dynamically relaxed., Many studies in recent years have attempted to separate cluster samples into two distinct populations: merging clusters vs clusters that are dynamically relaxed.661 The global properties of the latter are expected to be tightly correlated with mass. and therefore directly relatable to statistical predictions from cosmological models.," The global properties of the latter are expected to be tightly correlated with mass, and therefore directly relatable to statistical predictions from cosmological models."662 While there is evidence that some X-ray cluster properties may in fact be bimodal. e.g. with respect to the presence or absence of a cool core (?).. cosmological simulations have shown that even the most relaxed cool core clusters can still be accreting significant amounts of mass (?).. at a mean rate which is approximately linearly proportional to the mass of the system (?)..," While there is evidence that some X-ray cluster properties may in fact be bimodal, e.g. with respect to the presence or absence of a cool core \citep{Sanderson2009}, cosmological simulations have shown that even the most relaxed cool core clusters can still be accreting significant amounts of mass \citep{Poole2008}, at a mean rate which is approximately linearly proportional to the mass of the system \citep{McBride2009}."663 Generally. only the most massive of these accretion events (mass ratios <10 1) are expected to disturb the dense cluster core.," Generally, only the most massive of these accretion events (mass ratios $<10:1$ ) are expected to disturb the dense cluster core."664 Massive cool core clusters can thus quietly acerete a significant fraction of their total mass without large disturbances to their central ICM. which makes them arguably “cleaner” targets for studies of cluster galaxy evolution.," Massive cool core clusters can thus quietly accrete a significant fraction of their total mass without large disturbances to their central ICM, which makes them arguably “cleaner"" targets for studies of cluster galaxy evolution."665 The Local Cluster Substructure Survey (LoCuSS!)) key programme (?) was designed to probe dusty forming galaxies 1n the infall regions of a sample of 30 galaxy clusters ἐς~ 0.2) spanning a wide range of mass and merging histories.," The Local Cluster Substructure Survey ) key programme \citep{Smith2010}666 was designed to probe dusty star-forming galaxies in the infall regions of a sample of 30 galaxy clusters $z\sim0.2$ ) spanning a wide range of mass and merging histories."667 Abell 1835 is the most X-ray luminous cluster in. this sample (Zxpop=5.32+0.15x107 ergs/s (?))) and the host of a strong cool core (?)..," Abell 1835 is the most X-ray luminous cluster in this sample $L_{\rm668 X,bol}=5.32\pm 0.15 \times10^{45}$ ergs/s \citep{Zhang2008}) ) and the host of a strong cool core \citep{Peterson2001}."669 Its strong- and weak-lensing derived mass distributions (??).. and X-ray morphology (?) are typical of undisturbed nnon-merging) cool core clusters.," Its strong- and weak-lensing derived mass distributions \citep{Richard2010,Okabe2010}, and X-ray morphology \citep{Smith2005} are typical of undisturbed non-merging) cool core clusters."670 For example. the low substructure mass fraction in its core (Au= 0.13+0.01) suggests that it has grown in mass by <10% in the previous few Gyr (?)..," For example, the low substructure mass fraction in its core $f_{\rm sub}=0.13\pm0.01$ ) suggests that it has grown in mass by $<10\%$ in the previous few Gyr \citep{Smith2008}."671 In this Letter. we focus on asample of IOO0um-selectec spectroscopically confirmed cluster galaxies in 11835.," In this Letter, we focus on asample of $100\mu{\rm m}$ -selected spectroscopically confirmed cluster galaxies in 1835."672 We use the Photodetector Array Camera and Spectrometer (PACS) (?) on the satellite (2) to measure far infrared fluxes and from these we estimate total infrared luminosities for these objects., We use the Photodetector Array Camera and Spectrometer (PACS) \citep{Poglitsch2010} on the satellite \citep{Pilbratt2010} to measure far infrared fluxes and from these we estimate total infrared luminosities for these objects.673 We then use these data to examine the spatial and velocity distribution of obscured star formation sites i1 Abell 1835., We then use these data to examine the spatial and velocity distribution of obscured star formation sites in Abell 1835.674 We assume Hy= 70km s7!.Q4;=0.3.Q40.7.," We assume $H_0= 70$ km $^{-1}, \Omega_M =0.3,\Omega_\Lambda=0.7$."675 All virial masses and radit are stated with respect to the virial overdensity. Avia(<=0.25)121 (2)..," All virial masses and radii are stated with respect to the virial overdensity, $\Delta_{\rm676 virial}(z=0.25) =121$ \citep{Bryan1998}."677 PACS data were obtained in scan map mode on December 24th 2009. at both 100 and 160m. The details of the observing mode. data reduction procedure and source extraction and photometry can be found in (?)..," PACS data were obtained in scan map mode on December 24th 2009, at both 100 and $\mu$ m. The details of the observing mode, data reduction procedure and source extraction and photometry can be found in \citep{Smith2010}. ."678 The 90% completeness flux, The $90\%$ completeness flux679We now compare these model calcul1ος with the available abundance imeasurenieuts in he extremely metal-poor DLA reported by Cookeetal.(20101).filled circles)...,We now compare these model calculations with the available abundance measurements in the extremely metal-poor DLA reported by \citet[filled circles]{coo10b}.680 This DLA was originally identified in the Sloan Digita Sky Survey (SDSS) spectrmm of the QSO 0015., This DLA was originally identified in the Sloan Digital Sky Survey (SDSS) spectrum of the QSO $-$ 0918.681" Follow-up high resolution spectroscopy tudicated tan,=23100972. logN(HD. fem>=20.55X 0.1. Fe/H] =3.0L aud a pronounced carbon euliancemoent =|1.53."," Follow-up high resolution spectroscopy indicated $z_{\rm abs}=2.3400972$, $\log N({\rm HI})$ $/{\rm cm}^{-2}=20.55 \pm 0.1$ , [Fe/H] $\simeq -3.04$, and a pronounced carbon enhancement $\simeq +1.53$."682 Tn order to discuss the detailed. explosion mechauisui of the supernova. if is necessary to obtain the elenatal abundances of iron-peak clemenuts.," In order to discuss the detailed explosion mechanism of the supernova, it is necessary to obtain the elemental abundances of iron-peak elements."683 For the C-rich E)LA. because of the low metallicity. it is mupossible to detect heavier clements than S except for Fe.," For the C-rich DLA, because of the low metallicity, it is impossible to detect heavier elements than S except for Fe."684 Towever. the abundance proflius approach can be used for ¢ther absorption svstenis with peculiar anndancee patterns. because such systems lave prestunably been eunched bv only a sinall nuniber of superuovae.," However, the abundance profiling approach can be used for other absorption systems with peculiar abundance patterns, because such systems have presumably been enriched by only a small number of supernovae."685 Iu Figure 1 we overplot (open circles) the clement abundances recently reported by Cooleetal.(201Qa} for the tan.= DLA in front of the eravitationally lensed quasar UMG637A. This DLA. which has logeiN(HID/cn?=20.7£0.1 aud |Fe/H] =1.56+(0.03. exhibifs a different abundance pattern— from other DLAs with simular mctallicities aud also from the average popul:iion of Galactic metal-poor stars.," In Figure \ref{fig:dla} we overplot (open circles) the element abundances recently reported by \citet{coo10} for the $z_{\rm abs}=1.62650$ DLA in front of the gravitationally lensed quasar UM637A. This DLA, which has $\log N({\rm HI})/{\rm cm}^{-2}=20.7 \pm 0.1$ and [Fe/H] $= -1.56 \pm 0.03$, exhibits a different abundance pattern from other DLAs with similar metallicities and also from the average population of Galactic metal-poor stars."686 Iu particular. Ti. Ni. aud Zu are deficient relative to Fe (see Fie.," In particular, Ti, Ni, and Zn are deficient relative to Fe (see Fig."687 11 aud 12 in €OOLCal.2010233)., 11 and 12 in \citealt{coo10}) ).688 We show that this abundance pattern can al«) bC explained with faint supernovae., We show that this abundance pattern can also be explained with faint supernovae.689 The normal {(Si.S ratios are consistent with the supernova scenario.," The normal [(Si,S)/Fe] ratios are consistent with the supernova scenario."690 The ow Al abundance stronely sugeests that the euricluneut source is primordial supernovac., The low Al abundance strongly suggests that the enrichment source is primordial supernovae.691 The observed Zu/Fe ratio is more consistent with the faint SN model {solic ine) than the faint TIN model (dashed line) or Ol‘aint superuovae without imixing-fallback (|Zu/Ffon L.7 in WOG)., The observed [Zn/Fe] ratio is more consistent with the faint SN model (solid line) than the faint HN model (dashed line) or non-faint supernovae without mixing-fallback ([Zn/Fe] $\sim -1.7$ in K06).692 This is also supported by the observe Ni/Fo] |, This is also supported by the observed [Ni/Fe]. [693Cr/Fo] is consistent with both superLOVa nodels.,Cr/Fe] is consistent with both supernova models.694 We should note that there is a problem in tljo Ti imcleosvuthesis vields (I&06). aud we do uot incude Ti in the abundance profiling.," We should note that there is a problem in the Ti nucleosynthesis yields (K06), and we do not include Ti in the abundance profiling."695 At this metallicity. however. here is no information on C cuhancement in the DLA observations since € line is saturated.," At this metallicity, however, there is no information on C enhancement in the DLA observations since C line is saturated."696 Therefore. it is unucertain how eficieut the müxiug-fallback is.," Therefore, it is uncertain how efficient the mixing-fallback is."697 In act. he observed abundances can be explained witli micejuni uixiue-fallback of non-füut superuovae.," In fact, the observed abundances can be explained with medium mixing-fallback of non-faint supernovae."698 Iu this case. the ejected Fe mass is larger. which may be consisteut with he observed inetallicitv. and thus [C/Fe] is lower han ‘or the C-vich DLA.," In this case, the ejected Fe mass is larger, which may be consistent with the observed metallicity, and thus [C/Fe] is lower than for the C-rich DLA."699 Tn Figure d the dotted Lue is for the nucleosvutLOSIS vields of the PISN of a 1703Εν star. which are taken Yolu Uineda&Nomoto(2002).," In Figure \ref{fig:dla}700 the dotted line is for the nucleosynthesis yields of the PISN of a $170M_\odot$ star, which are taken from \citet{ume02}."701". The primordial stars with initial masses of ~110Στυλ, cuter into the clectron-positron pair mstabilitv region during the ceutral oxvecu-burning stages and contract quasi-dvnamically.", The primordial stars with initial masses of $\sim 140-270 M_\odot$ enter into the electron-positron pair instability region during the central oxygen-burning stages and contract quasi-dynamically.702 Then the ceutral temperature mereases. ceutral oxveen buiue takes place explosively. aud the generated nuclear energy is large enough to disrupt the stars completely without leaving compact remnants.," Then the central temperature increases, central oxygen burning takes place explosively, and the generated nuclear energy is large enough to disrupt the stars completely without leaving compact remnants."703 Compared with core-collapse supernovac. the abundance pattern of PISNe cau be stnunarized as follow.," Compared with core-collapse supernovae, the abundance pattern of PISNe can be summarized as follow."704The powerlaw fits show no indications that the quicscent lightcurve of iis levelling olf.,The powerlaw fits show no indications that the quiescent lightcurve of is levelling off.705 Thus. it is also possible that the neutron star temperature continues to decay further and that the core is cooler than suggested by the exponential decay fits and the 1980 dadetection.," Thus, it is also possible that the neutron star temperature continues to decay further and that the core is cooler than suggested by the exponential decay fits and the 1980 detection."706 The relatively slow decrease of mmight then rellect that the crust has a high conductivity. albeit lower than that of the neutron stars in aand29.," The relatively slow decrease of might then reflect that the crust has a high conductivity, albeit lower than that of the neutron stars in and."707. Further observations are thus required to determine whether the neutron star crust in hhas nearly cooled clown and to be able to draw firm conclusions on the crust and core properties., Further observations are thus required to determine whether the neutron star crust in has nearly cooled down and to be able to draw firm conclusions on the crust and core properties.708 This work was supported by the Netherlands Organisation for Scientific Research (NWO) and made use of the ppublic data archive., This work was supported by the Netherlands Organisation for Scientific Research (NWO) and made use of the public data archive.709 We acknowledge DDPLN. Gehrels and the pplanning team for their help in carrying out the ToO campaign., We acknowledge PI N. Gehrels and the planning team for their help in carrying out the ToO campaign.710 IZMC' was supported by NASA through the Chandra Fellowship Program., EMC was supported by NASA through the Chandra Fellowship Program.711 AITW. PSR and. WSW acknowledge the United. States Ollice of Naval Rescarch.," MTW, PSR and KSW acknowledge the United States Office of Naval Research."712 aancd acknowledge support from Chandra erant GOS-9045X., and acknowledge support from Chandra grant GO8-9045X.713"by c and ες respectively. are related by where the complex quantity g is the ""reduced: shear"". with its complex conjugate denoted by g.","by $\epsilon$ and $\epsilon_s$ respectively, are related by where the complex quantity $g$ is the “reduced shear”, with its complex conjugate denoted by $g^*$."714 Phe reduced shear is then related. to the (complex) shear. 5. (the tidal eravitational field). and the convergence. s. by Equation implies that in the ideal case of a perfectly circular background galaxy with r = 1 and c;=0. the lensed ellipticitv €=g.," The reduced shear is then related to the (complex) shear, $\gamma$ (the tidal gravitational field), and the convergence, $\kappa$, by Equation implies that in the ideal case of a perfectly circular background galaxy with r = 1 and $\epsilon_s = 0$, the lensed ellipticity $\epsilon = g$."715" In the above. & is proportional to the projected. mass density of the lens. X: where is the ""critical density and 1). Oy and D; are the angular diameter distances between observer source. observer lens and lens ποιος respectively."," In the above, $\kappa$ is proportional to the projected mass density of the lens, $\Sigma$: where is the “critical density” and $D_s$, $D_d$ and $D_{ds}$ are the angular diameter distances between observer – source, observer – lens and lens – source respectively."716 In this notation. the WL regime is where X««X. Le. 8««I.," In this notation, the WL regime is where $\Sigma << \Sigma_{\text{cr}}$, i.e. $\kappa <<1$."717 The lensing dellection potential. c. is related τος and HB by a set of partial dillerential equations (e.g. ?)): where and the indices following the comma denote partial derivatives with respect to the components of the position vector 8.," The lensing deflection potential, $\psi$, is related to $\gamma$ and $\kappa$ by a set of partial differential equations (e.g. \citealt{Bartelmann_Schneider_2001}) ): where and the indices following the comma denote partial derivatives with respect to the components of the position vector $\mathbf{\theta}$."718 A secondary effect. is a slight modification in the observed. number density of background. galaxies. Πρι. due to magnification. changing the apparent [uxes of galaxies. as well as changing the apparent area of sky. in which they are observed.," A secondary effect is a slight modification in the observed number density of background galaxies, $n_{\rm lensed}$, due to magnification changing the apparent fluxes of galaxies, as well as changing the apparent area of sky in which they are observed."719" Ehe net effect depends on the slope ofthe number counts of the galaxies 67). with the lensed and unlensed number counts being related by (2): where the magnification yr is given hy Since 3s0.5 for the faint distant. galaxies tvpically used. in. WL analyses. this results. in. “number count depletion"" (7)... with the observed background. galaxy density decreasing with decreasing distance from the cluster centre in the WL regime."," The net effect depends on the slope of the number counts of the galaxies $\beta$ ), with the lensed and unlensed number counts being related by \citep{Canizares_1982}: where the magnification $\mu$ is given by Since $\beta\approx 0.5$ for the faint distant galaxies typically used in WL analyses, this results in “number count depletion"" \citep{Broadhurst_et_al_1995}, with the observed background galaxy density decreasing with decreasing distance from the cluster centre in the WL regime."720 We extract simulated. galaxy clusters from the Alillenniun Simulation (hereafter MS: 2)). a laree cosmological N-body simulation that follows 2160? dark matter particles from 2=127 to 2=0 in a periodic box of 5005 MMpeO on a side.," We extract simulated galaxy clusters from the Millennium Simulation (hereafter MS; \citealt{Springel_et_al_2005}) ), a large cosmological N-body simulation that follows $2160^3$ dark matter particles from $z=127$ to $z=0$ in a periodic box of $500\,h^{-1}$ Mpc on a side."721 The cosmology adopted. for the MS. which we also adopt for our analysis. is a flat ACDAL moclel with f=0.73. Ox;=0.25. ancl a power spectrum normalisation on a scale of Sh‘Alpe ol ox=0.9 (οι. the rms linear mass Iluctuation in a sphere of radius sf Ape extrapolated to present-day).," The cosmology adopted for the MS, which we also adopt for our analysis, is a flat $\Lambda$ CDM model with $h=0.73$, $\Omega_{\rm M} =0.25$, and a power spectrum normalisation on a scale of $8 h^{-1}$ Mpc of $\sigma_8 = 0.9$ (i.e., the rms linear mass fluctuation in a sphere of radius $8 h^{-1}$ Mpc extrapolated to present-day)."722 “Phese parameters are. consistent with the latest measurements of temperature ancl polarization anisotropies in the cosmic microwave hackeround (CALB) with the (WALAD: ?2)). although the value of ex adopted for the MS is larger than the maximum likelihood CAIB value by z260.," These parameters are consistent with the latest measurements of temperature and polarization anisotropies in the cosmic microwave background (CMB) with the (WMAP; \citealt{Komatsu_et_al_2011}) ), although the value of $\sigma_8$ adopted for the MS is larger than the maximum likelihood CMB value by $\approx 2\sigma$."723 The larger value of ox means that the ALS will have more massive clusters than a universe with the WALAD T-vear cosmology., The larger value of $\sigma_8$ means that the MS will have more massive clusters than a universe with the WMAP 7-year cosmology.724 Llowever. as we are interested in the lensing signal of clusters (specifically how well WL can recover their mass anc concentration) and not their abundances. this discrepaney will not alfect the validity of our results.," However, as we are interested in the lensing signal of clusters (specifically how well WL can recover their mass and concentration) and not their abundances, this discrepancy will not affect the validity of our results."725 From the snapshot at z0.2. a typical redshift. of observed WL clusters. we select all of the simulated clusters with masses Mou20HAI. for analysis.," From the snapshot at $z \simeq 0.2$, a typical redshift of observed WL clusters, we select all of the simulated clusters with masses $M_{200} \ge 10^{14} M_\odot$ for analysis."726" Note that Aon is defined as the mass within a radius (705) that eneloses a mean density equal to 200 times the critical density of the universe in our adopted cosmology. this density. is Pot=LIGAOMAL,Alpe? at our chosen redshift."," Note that $M_{200}$ is defined as the mass within a radius $r_{200}$ ) that encloses a mean density equal to 200 times the critical density of the universe — in our adopted cosmology, this density is $\rho_{\text{crit}} = 1.76 \cdot 10^{11} M_\odot / \text{Mpc}^3$ at our chosen redshift."727" Fron hore on. we will refer to these masses obtained cirecthy from the AIS particle cistribution without any profile fitting as ""true mass. denoted by Aii."," From here on, we will refer to these masses obtained directly from the MS particle distribution without any profile fitting as `true' mass, denoted by $\mmill$."728 A large sample of 2678 simulate cluster haloes satisfies this criterion. which will allow us to robustIy quantify the mean trends and scatter in the derivec niss-concentration relationship.," A large sample of 2678 simulated cluster haloes satisfies this criterion, which will allow us to robustly quantify the mean trends and scatter in the derived mass-concentration relationship."729 We extract all of the dark matter particles within a 105+ AIAIpe (comoving) box. centered. on the mos bound: particle. for. producing WL maps of cach cluster.," We extract all of the dark matter particles within a $10\,h^{-1}$ Mpc (comoving) box, centered on the most bound particle, for producing WL maps of each cluster."730 For reference. rego is typically =2 Alpe for the mos massive clusters in our sample (ancl smaller for lower mass clusters).," For reference, $r_{200}$ is typically $\approx 2$ Mpc for the most massive clusters in our sample (and smaller for lower mass clusters)."731 Thus. our analysis includes only thelocal environment around the clusters.," Thus, our analysis includes only the environment around the clusters."732" This was deliberate. so that we can isolate the ellects of cluster triaxiality. substructure and. connecting filaments (Le. ""correlated? signals) from. uncorrelatect line-of-sight structures at much larger distances."," This was deliberate, so that we can isolate the effects of cluster triaxiality, substructure and connecting filaments (i.e., “correlated” signals) from uncorrelated line-of-sight structures at much larger distances."733 While these correlated: structures can in general extend bevond a distance of 5 f+ from the cluster centre. we verify in appendix Bl (Fie. D1))," While these correlated structures can in general extend beyond a distance of 5 $h^{-1}$ from the cluster centre, we verify in appendix \ref{sec:testlength} (Fig. \ref{fig:boxsizetest}) )"734 that. in agreement with the findings by 2.. our results are robust o increasing the linc-of-xight integration length by a factor of five to 50h comoving Alpe.," that, in agreement with the findings by \citet{Becker_Kravtsov_2011}, our results are robust to increasing the line-of-sight integration length by a factor of five to $50\,h^{-1}$ comoving Mpc."735 Much. longer. integration eneths. which would capture uncorrelated larec scale structure more fully. would in principle be desirable for the purpose of improving the prediction accuracy. but. require 10 use of rav-tracing algorithms which is bevond the scope of this work.," Much longer integration lengths, which would capture uncorrelated large scale structure more fully, would in principle be desirable for the purpose of improving the prediction accuracy, but require the use of ray-tracing algorithms which is beyond the scope of this work."736 In a complementary studs. 2? have used rav- of the MS to examine the ellect of such uncorrelated structures on the derived. masses and: concentrations of analytic clusters and found that it contributes to scatter. but," In a complementary study, \citet{Hoekstra_et_al_2011} have used ray-tracing of the MS to examine the effect of such uncorrelated structures on the derived masses and concentrations of analytic clusters and found that it contributes to scatter, but"737broad emission roffesults)) aud on the interpretation of this emission rofDiscussion)).,broad emission \\ref{Results}) ) and on the interpretation of this emission \\ref{Discussion}) ).738 From these fits to the average spectra. we nmnoeasure the fractional contribution of broad ciaission to the overall ciuission line flux. the kinematics of the broad component. and the Lue ratios of the narrow line compoucuts.," From these fits to the average spectra, we measure the fractional contribution of broad emission to the overall emission line flux, the kinematics of the broad component, and the line ratios of the narrow line components."739 The significance of these irieasuremoents are quantified in two wavs., The significance of these measurements are quantified in two ways.740 First. for cach average spectrum. we re-created 100 spectra by randomly sampling (with replacement) aud combining the idividual coutzibutiug ealaxy spectra.," First, for each average spectrum, we re-created 100 spectra by randomly sampling (with replacement) and combining the individual contributing galaxy spectra."741 The properties of the broad component were ineasured iu cach of the resulting LOO average spectra. viellius confidence intervals for all derived quantities.," The properties of the broad component were measured in each of the resulting 100 average spectra, yielding confidence intervals for all derived quantities."742 Second. the probability of false positive detections Pyare was tested by creating 1000. simulated spectra with the narrow line aud observational properties (noise. spectral resolution) characteristic of cach average spectrum.," Second, the probability of false positive detections $P_{false}$ was tested by creating 1000 simulated spectra with the narrow line and observational properties (noise, spectral resolution) characteristic of each average spectrum."743 Comparing the derived broad compoucuts in these spectra to those im the real SINS spectra. we estimate the rate of spurious detections of broad conrponeuts equal to or more pronünent (n huuinositv and EWIIMS than those in the actual data.," Comparing the derived broad components in these spectra to those in the real SINS spectra, we estimate the rate of spurious detections of broad components equal to or more prominent (in luminosity and FWHM) than those in the actual data."744 The results of this analvsis for the cdiffereut average spectra are preseuted in the followiug section., The results of this analysis for the different average spectra are presented in the following section.745" We find that the average SINS galaxw spectrum iucludes a significant amount of broad cmission (top xuel of Figure 1)). with v), 1.9 for a fit with oulv jurow lines (dashed magenta line in top right paucl of Figure 1)) aud ot l.5 and l.l for fits iucludiug a xoad line aud broad aaud wwines. [NH]respectively (red aud dotted ereeu lines in same auch Prose= and 1%))."," We find that the average SINS galaxy spectrum includes a significant amount of broad emission (top panel of Figure \ref{agn}) ), with $\chi^2_{dof}$ = 4.9 for a fit with only narrow lines (dashed magenta line in top right panel of Figure \ref{agn}) ) and $\chi^2_{dof}$ = 1.8 and 1.4 for fits including a broad line and broad and wings, respectively (red and dotted green lines in same panel; $P_{false} =$ and )."746 To cusure that this signature is not the result of the 1 known ACN included in this sample. we also create stacked spectra of the AGN and the rest of the sample.," To ensure that this signature is not the result of the 4 known AGN included in this sample, we also create stacked spectra of the AGN and the rest of the sample."747 While the broad ciission roni the ACN host svstenis alone is quite substantial. a colparison of the average SINS spectrum and the average non-AQGN spectrum illustrates that the broad enission in the average SINS προςπα is uot douinated w that coming from the | AGN.," While the broad emission from the AGN host systems alone is quite substantial, a comparison of the average SINS spectrum and the average non-AGN spectrum illustrates that the broad emission in the average SINS spectrum is not dominated by that coming from the 4 AGN."748 These results are sununarized in Table , These results are summarized in Table \ref{tab:results}.749We test the dependence1 of the presence of broad Cluission on galaxy properties by dividing the sauple iuto three stellar 1iass bius. using the results from the spectral energv distribution (SED) fittine of for the oof our areets for which sufficient broadband. data exist.," We test the dependence of the presence of broad emission on galaxy properties by dividing the sample into three stellar mass bins, using the results from the spectral energy distribution (SED) fitting of for the of our targets for which sufficient broadband data exist."750 The average spectra for these three bins show au Increasing presence of a broad component with stellar nass (Figure 2: Prone= I1. respectively).," The average spectra for these three bins show an increasing presence of a broad component with stellar mass (Figure \ref{mstar}; $P_{false} =$ , respectively)."751 Uowever. the spectrum of the highest amass biu is sienificautly affected by the contiubution of the | xeviouslv known ACN our sample. which all fall iuto this biu.," However, the spectrum of the highest mass bin is significantly affected by the contribution of the 4 previously known AGN in our sample, which all fall into this bin."752 The nuclear cussion in these svsteiis can jas the results of SED fitting towards larger masses. so we confirm thei high masses with the dvuamical mass neasurclents made by(2009):: tn all cases. he dvnaical masses of these galaxies are consisteut with the stellar masses used here and remain among the highest in the SINS sample.," The nuclear emission in these systems can bias the results of SED fitting towards larger masses, so we confirm their high masses with the dynamical mass measurements made by; in all cases, the dynamical masses of these galaxies are consistent with the stellar masses used here and remain among the highest in the SINS sample."753 Nevertheless. we confirm that a Qveaker) broad compoucut is also present iu the nuon-active galaxies iu this bin (ereen Lue in Figure 2)). with Pius=key ," Nevertheless, we confirm that a (weaker) broad component is also present in the non-active galaxies in this bin (green line in Figure \ref{mstar}) ), with $P_{false} =$."754"Several other v properties of ealaxies. including SER. size, stellar age. and imetalliitv. have well established correlations with stellar mass in eealaxies»)."," Several other key properties of galaxies, including SFR, size, stellar age, and metallicity, have well established correlations with stellar mass in galaxies."755. Both the ielation and the mass-nctallicity relation are apparcut in Fieure 2.. via the increasing narrow Thuninosity aud increasing rratio with stellar mass. respectively.," Both the relation and the mass-metallicity relation are apparent in Figure \ref{mstar}, via the increasing narrow luminosity and increasing ratio with stellar mass, respectively."756" Since the ratio remains well below levels expected of shock[NU heating or ACN activity, this latter is most likely tracing variations in iuetallicity (sce below. as well as Buschkamp et al."," Since the ratio remains well below levels expected of shock heating or AGN activity, this latter is most likely tracing variations in metallicity (see below, as well as Buschkamp et al."757 iu prep)., in prep).758 To these established mass-depeudoeut properties im eealaxies. we now add the presence and strength of a broad component.," To these established mass-dependent properties in galaxies, we now add the presence and strength of a broad component."759 With the spatially resolved data. we cau also compare the iutegrated spectra from the central (Ro<3 kpc) regions of eealaxies to those from extended (Ro=5/15 kpc) regions. iu order to determine what regions in these systelus are eeneratimg broad emission.," With the spatially resolved data, we can also compare the integrated spectra from the central $R~<~3$ kpc) regions of galaxies to those from extended $R~=~3-15$ kpc) regions, in order to determine what regions in these systems are generating broad emission."760 For this analysis. we use galaxies frou the intermediate and high mass bius of Fieure 2. in which the intensity distribution of the stellar ουή. defines a clear center of the system (totalling 6 svstcms).," For this analysis, we use galaxies from the intermediate and high mass bins of Figure \ref{mstar} in which the intensity distribution of the stellar continuum defines a clear center of the system (totalling $6$ systems)."761 The average spectra of the ceutral aud extended regions of these svstceius are slow in Figure 3.. normalized to the spatial area over which the spectra were extracted.," The average spectra of the central and extended regions of these systems are shown in Figure \ref{cenext}, normalized to the spatial area over which the spectra were extracted."762 Tn these spectra. a broad componcut is preferred bv the best-fitting models: however. the significance of this result is low.," In these spectra, a broad component is preferred by the best-fitting models; however, the significance of this result is low."763 The detection of broad cutission iu galaxy centers at (vith Pyare=8% )) is somewhat more robust than that in the extended regious (Prats 11). in which the vost fit broad compoucut is very shallow.," The detection of broad emission in galaxy centers at (with $P_{false} = 8$ ) is somewhat more robust than that in the extended regions $P_{false} = 14$ ), in which the best fit broad component is very shallow."764 If real. the woad feature in the extended regious accounts for a comparable fraction of the total huninosity to that iu the ceutral regions (Table 1)).," If real, the broad feature in the extended regions accounts for a comparable fraction of the total luminosity to that in the central regions (Table \ref{tab:results}) )."765 Tests of simulated ealaxies indicate that such a broad liue iu he extended regious cannot be reproduced by a nuclear xoimt-source of broad. cussion(ie. ACN)broadenued by he PSF., Tests of simulated galaxies indicate that such a broad line in the extended regions cannot be reproduced by a nuclear point-source of broad emission (i.e. AGN) broadened by the PSF.766 Tn the spectra shown iu Figure 3.. we note that the average ratio is comparable iu the central and exteuded regions. despite the difference in broad ciissiou compoucut. i support of the above interpretation of the ffeature as primarily reflecting the metallicity of these systenis.," In the spectra shown in Figure \ref{cenext}, we note that the average ratio is comparable in the central and extended regions, despite the difference in broad emission component, in support of the above interpretation of the feature as primarily reflecting the metallicity of these systems."767 However. this should not be interpreted as a lack of unctallicity eracieut in these galaxies. since the spectra shown here are averaged over a nuuber of galaxies. im each of which the extended region is spatially iutegrated over a large range in radii.," However, this should not be interpreted as a lack of metallicity gradient in these galaxies, since the spectra shown here are averaged over a number of galaxies, in each of which the extended region is spatially integrated over a large range in radii."768 Detailed studies of metallicity eradieunts within aud between individual SINS svstenis will be preseuted in Buschkamp et al. (, Detailed studies of metallicity gradients within and between individual SINS systems will be presented in Buschkamp et al. (769n prep).,in prep).770 The low luninosity broad cussion seen in our SINS, The low luminosity broad emission seen in our SINS771strength lead to estimates of  6 mG. This is the strongest magnetic field. calculated for the cores in this sample. but is smaller than the field strengths. calculated from. the maser emission (Vlemminegsetal.2006).,"strength lead to estimates of $\sim$ 6 mG. This is the strongest magnetic field calculated for the cores in this sample, but is smaller than the field strengths calculated from the maser emission \citep{Vlemmings06}."772. This suggests hat the magnetic field is weaker on larger scales. becoming stronger on smaller scales. closer in to the protostar.," This suggests that the magnetic field is weaker on larger scales, becoming stronger on smaller scales, closer in to the protostar."773 The data presented. here. in conjunction with the findings of Dartkiewiczctal.(2005). and Vlemamingsctal.(2006). ead to a magnetic field that. in the plane of the sky. is northeast-southwest in direction. with the magnetic field in he northeast pointing away from us. and in the southwest »ointing toward us.," The data presented here, in conjunction with the findings of \citet{Bartkiewicz} and \citet{Vlemmings06}, lead to a magnetic field that, in the plane of the sky, is northeast-southwest in direction, with the magnetic field in the northeast pointing away from us, and in the southwest pointing toward us."774 DIHA21I(O00D is part of the star forming complex W75. located in the Cygnus X region. at a distance of ~ 3 kpe (Campbell 1982).," DR21(OH) is part of the star forming complex W75, located in the Cygnus X region, at a distance of $\sim$ 3 kpc \citep{campbell}."775.. Hs relative close proximity has aided. in the fact that it is one of the most studied star forming regions in the Galaxy., Its relative close proximity has aided in the fact that it is one of the most studied star forming regions in the Galaxy.776 The DBR21 region is ~ to the south of DR21COLD) ancl contains a cluster of late vpe O stars. and one of the most intense outllows known.," The DR21 region is $\sim$ to the south of DR21(OH) and contains a cluster of late type O stars, and one of the most intense outflows known."777 (OID | also known as W758 or W75S(OLD) — unlike . consists of voung stars in the process of forming. anc as such ollers a glimpse of an earlier stage of evolution of star formation.," DR21(OH) – also known as W75S or W75S(OH) – unlike DR21, consists of young stars in the process of forming, and as such offers a glimpse of an earlier stage of evolution of star formation."778 Previous continuum studies have shown tha (OLD is made up of three compact continuum sources. (OLD) Main. DR21(O01DS. and. DIC2ICOLDW. all of which seem to be actively forming stars (Mangum.Wootten&Mundy 1991).," Previous continuum studies have shown that DR21(OH) is made up of three compact continuum sources, DR21(OH) Main, DR21(OH)S and DR21(OH)W, all of which seem to be actively forming stars \citep{mangum91}."779 Higher resolution studies confirmed. tha DRTOLDMain is in fact composed of two smaller cores. AIME and MAIZ (Woodsetal.1989).," Higher resolution studies confirmed that DR21(OH)Main is in fact composed of two smaller cores, MM1 and MM2 \citep{woody89}."780 The magnetic field. of ΟΡΙΟ) is reasonably wel studied. and there are Zeeman estimates of the line of sieh magnetic field strength for MALL ancl MM2 (Crutcherctal. 1999).," The magnetic field of DR21(OH) is reasonably well studied, and there are Zeeman estimates of the line of sight magnetic field strength for MM1 and MM2 \citep{crutcher99}."781. Laietal.(2003b) usec BIALA to eain information on the magnetic field via polarimetry of both dust and CO. in order to map the magnetic field morphology in the plane of the sky.," \citet{lai04} used BIMA to gain information on the magnetic field via polarimetry of both dust and CO, in order to map the magnetic field morphology in the plane of the sky."782 Comparisons of ion/neutral linowidths have also been used to establish a three-cdimensional impression of the magnetic field. (Laietal.2003a)., Comparisons of ion/neutral linewidths have also been used to establish a three-dimensional impression of the magnetic field \citep{lai03}.783. The morphology. of the field at the resolution sampled by BIALA indicates that the magnetic field is ordered. thus implving a strong field.," The morphology of the field at the resolution sampled by BIMA indicates that the magnetic field is ordered, thus implying a strong field."784 The magnetic field is estimated to have a strength of ~ 0.4 mG with an inclination of 86° to the line of sight. ancl a position angle of 105 in the plane of the sky (Laietal.2003a).," The magnetic field is estimated to have a strength of $\sim$ 0.4 mG with an inclination of $^{\circ}$ to the line of sight, and a position angle of $^{\circ}$ in the plane of the sky \citep{lai03}."785. The BIALA data only measure the polarisecl dust. emission in patches. and do not reveal the magnetic field morphology throughout the whole region.," The BIMA data only measure the polarised dust emission in patches, and do not reveal the magnetic field morphology throughout the whole region."786 The data presented here (Eig Lbb) include our observed data as well as some archival data for. ΟΣΟΙ). which were observed on 2002 October 2 (Vallée&Fiege2006).. which effectively doubles the time on source for this target.," The data presented here (Fig \ref{fig1}b b) include our observed data as well as some archival data for DR21(OH), which were observed on 2002 October 2 \citep{vallee}, which effectively doubles the time on source for this target."787 The archival data were subjected to our method. of data reduction as detailed: earlier., The archival data were subjected to our method of data reduction as detailed earlier.788 The data reveal that the DIH21(O01I) region is composed. (at this resolution) of one main core DBR21COLD) Main., The data reveal that the DR21(OH) region is composed (at this resolution) of one main core DR21(OH) Main.789 There are also two fainter cores close to the main core. one to the south-west. DR21CO0LDW and one due south. DI21(01DS.. Mangum.Wootten&Alundy(1992) identified another core. DIR21(COLIDN. in NIL; emission. although in the SCUBA data presented here it appears to have the morpholoev of a ridge. extending northwards from the main core.," There are also two fainter cores close to the main core, one to the south-west, DR21(OH)W and one due south, DR21(OH)S. \citet{mangum92} identified another core, DR21(OH)N, in $_{3}$ emission, although in the SCUBA data presented here it appears to have the morphology of a ridge extending northwards from the main core."790 “Phere are also molecular outIlows associated with the main core cmanating in an west direction (Laiοἱal.2003b)., There are also molecular outflows associated with the main core emanating in an east-west direction \citep{lai04}.791. Our calculated mass of the Main core is somewhat larger than those of Mangum.Wootten&Mundy(1991).. but they calculated the individual masses of MMI and AIAI2 using OVRO interferometer measurements of the dust emission.," Our calculated mass of the Main core is somewhat larger than those of \citet{mangum91}, but they calculated the individual masses of MM1 and MM2 using OVRO interferometer measurements of the dust emission."792 It is to be expected that our calculations reveal higher masses. eiven that they are based on single dish measurements and so include more diffuse dust on larger scales.," It is to be expected that our calculations reveal higher masses, given that they are based on single dish measurements and so include more diffuse dust on larger scales."793 Comparisons of our derived total masses with those of Vallée&Fiege(2006) reveal our masses are much higher than their estimates. however they have assumed a temperature of 100 Ix for both the Main and northern source.," Comparisons of our derived total masses with those of \citet{vallee} reveal our masses are much higher than their estimates, however they have assumed a temperature of 100 K for both the Main and northern source."794 We also use larger apertures for our calculations., We also use larger apertures for our calculations.795 Once these two factors have been taken into account. the mass estimates are consistent.," Once these two factors have been taken into account, the mass estimates are consistent."796 The polarimetry indicates that the magnetic field. is ordered. across the ridge. parallel to the outllow axis.," The polarimetry indicates that the magnetic field is ordered across the ridge, parallel to the outflow axis."797 The percentage polarisation also remains stable (at ~ 3%)) across the ridge., The percentage polarisation also remains stable (at $\sim$ ) across the ridge.798 Across the main core the percentage polarisation drops. most apparently to the northeast of the core. coincident with MMMI.," Across the main core the percentage polarisation drops, most apparently to the northeast of the core, coincident with MM1."799 This is observed. in numerous other cores (for.example.—Chirvsostomouetal.2002:Matthews&Wilson2002:Davisetal.9000) and. could be due to the magnetic field. twisting within the JCMT beam. the grains becoming more spherical in regions of high density. or the grains being less ellicientlv aligned in regions of high density.," This is observed in numerous other cores \citep[for example, ][]{ant,brenda,chris} and could be due to the magnetic field twisting within the JCMT beam, the grains becoming more spherical in regions of high density, or the grains being less efficiently aligned in regions of high density."800 To the south of the main core. the vectors. across DIA21(OLDS are more dispersed in position angle. such that to the southwest of DI2ICOLDS. the vectors have à position angle of ~ 135. changing to ~ 90 northeast of the source.," To the south of the main core, the vectors across DR21(OH)S are more dispersed in position angle, such that to the southwest of DR21(OH)S, the vectors have a position angle of $\sim$ $^{\circ}$, changing to $\sim$ $^{\circ}$ northeast of the source."801 The vectors across the south-western core are  907. the same as across the ridge.," The vectors across the south-western core are $\sim$ $^{\circ}$ , the same as across the ridge."802 The polarimetry data. in general agree with the lindings of Vallée& Fiege.(2006)... however we see a smaller dispersion in position angles throughout the region. which mav be due to a combination of higher signal-to-noise observations and a more careful data reduction.," The polarimetry data, in general agree with the findings of \citet{vallee}, however we see a smaller dispersion in position angles throughout the region, which may be due to a combination of higher signal-to-noise observations and a more careful data reduction."803 The rolarumetry agree with the findings of Laietal.(2003b).. with the polarisation nulls coincident with AIALL. where Lai et al.," The polarimetry agree with the findings of \citet{lai04}, with the polarisation nulls coincident with MM1, where Lai et al."804 found very Little polarisation. from. dust., found very little polarisation from dust.805 The overall change in direction of the magnetic field. across DRO)Main also agrees with the BIAIA cata both 10 dust and the CO polarimetry — indicating that the MZuagnetie field stavs ordered on both large ancl small scales., The overall change in direction of the magnetic field across DR21(OH)Main also agrees with the BIMA data – both the dust and the CO polarimetry – indicating that the magnetic field stays ordered on both large and small scales.806 Phe observed polarisation nulls across MMI may indicate wt either MMI has a twisted magnetic field in comparison Oo AMAI2. or that AMIAIL is more centrally condensed than AMAI2.," The observed polarisation nulls across MM1 may indicate that either MM1 has a twisted magnetic field in comparison to MM2, or that MM1 is more centrally condensed than MM2."807 Both of these would be consistent with MMI being 1e more evolved of the two cores., Both of these would be consistent with MM1 being the more evolved of the two cores.808 The polarimetry across the ridge (northern source) of rw DIR21COLD) region indicate an ordered. field. which in --πο implies a strong field.," The polarimetry across the ridge (northern source) of the DR21(OH) region indicate an ordered field, which in itself implies a strong field."809 The field is (in the plane. of re sky) perpendicular to the north-south ridge. vielding 16 possibility that collapse has occurred. along the field ines.," The field is (in the plane of the sky) perpendicular to the north-south ridge, yielding the possibility that collapse has occurred along the field lines."810 CE calculations reveal the magnetic field. strength is of the order ~ 1 mC in the plane of the sky., CF calculations reveal the magnetic field strength is of the order $\sim$ 1 mG in the plane of the sky.811 This is comparable to the field strengths previously calculated. for he two sources inDIR2ICOLHDMain by Crutcheretal. and Laictal. (2003b).. which were  Lom in the plane," This is comparable to the field strengths previously calculated for the two sources inDR21(OH)Main by \citet{crutcher99} and \citet{lai04}, , which were $\sim$ 1 mG in the plane"812"1n order to recover the continuous density field for dark matter component. p,,,. we use triangular shaped. cloud (PSC) scheme (Ilockney. Eastwood LOSS) at cach time slop.","In order to recover the continuous density field for dark matter component, $\rho_{_{DM}}$, we use triangular shaped cloud (TSC) scheme (Hockney Eastwood 1988) at each time step."813 The gravitational potential is computed by solving Poisson's equation., The gravitational potential is computed by solving Poisson's equation.814" As two components are considered. gas ancl dark matter. the source in Poisson's equation is the total density contrast: where 6,=a,|4oa1 when à,=PoiPy] and δρρου1."," As two components are considered, gas and dark matter, the source in Poisson's equation is the total density contrast: where $\delta_{_{T}}=\delta_{b}+\delta_{_{DM}}+1$ when $\delta_{b}=\rho_b/\rho_{_{B}} -1$ and $\delta_{_{DM}}=\rho_{_{DM}}/\rho_{_{B}}-1$."815" Poisson's equation (22)) is solved. using Fast Fourier ""Transform. (EET) methods (Press et al.", Poisson's equation \ref{poisson}) ) is solved using Fast Fourier Transform (FFT) methods (Press et al.816 1996)., 1996).817 The FET is used as follows: In order to solve numerically the Eq (15)) and Eqs (16- 17)). we need to choose a time step.," The FFT is used as follows: In order to solve numerically the Eq \ref{runge}) ) and Eqs \ref{dm1}- \ref{dm2}) ), we need to choose a time step."818 Phe numerical stability of the methods used to integrate these equations imposes several criteria on the time step., The numerical stability of the methods used to integrate these equations imposes several criteria on the time step.819 At cach numerical iteration. we compute several time steps given by the dillerent. stability conditions.," At each numerical iteration, we compute several time steps given by the different stability conditions."820 Phe most restrictive of all of them is selected to advance the gaseous and dark matter components., The most restrictive of all of them is selected to advance the gaseous and dark matter components.821" The time step criteria we consider are the following: The elobal time step is defined as the most stringent of all the previous time steps: At the beginning of the cosmological simulations. Af, is the dominant time criteria but Aves ancl 2Mpa quickly take over."," The time step criteria we consider are the following: The global time step is defined as the most stringent of all the previous time steps: At the beginning of the cosmological simulations, $\Delta t_{e} $ is the dominant time criteria but $\Delta t_C$ and $\Delta t_{DM}$ quickly take over."822" The fundamental idea behind the AMI, technique is to overcome the lack of resolution associated with the fix grid Eulerian description.", The fundamental idea behind the AMR technique is to overcome the lack of resolution associated with the fix grid Eulerian description.823 The basic idea is simple., The basic idea is simple.824 Regions in the original computational domain in which improved resolution is required. are selected: according to some criteria. (see Sec., Regions in the original computational domain in which improved resolution is required are selected according to some criteria (see Sec.825 3.1)., 3.1).826 Fhese new computational domains. which we call orpalehes. are remapped with a higher number of cells and therefore with better resolution.," These new computational domains, which we call or, are remapped with a higher number of cells and therefore with better resolution."827 The values of the iferent quantities defined on the child grids are obtained by interpolating from thegrid., The values of the different quantities defined on the child grids are obtained by interpolating from the.828 Once the child grids are built. they can be evolved as an independent computational omain by using the same methods we have described in Sec 2.," Once the child grids are built, they can be evolved as an independent computational domain by using the same methods we have described in Sec 2."829 Although conceptually simple. there are severe technical complications concerning with the communication among 1e dillerent patehes and the boundary problems at different levels.," Although conceptually simple, there are severe technical complications concerning with the communication among the different patches and the boundary problems at different levels."830 Our implementation of the AMIR technique follows the one described in Berger Colella(1989)., Our implementation of the AMR technique follows the one described in Berger Colella(1989).831 The [first step in the construction of the hierarchy of patches is the coarse basic erid on which all the relevant quantities are known., The first step in the construction of the hierarchy of patches is the coarse basic grid on which all the relevant quantities are known.832 From this starting point. some criteria must be applied to decide which cells arerefénabíe.," From this starting point, some criteria must be applied to decide which cells are."833 These criteria are application dependent and may need. to. be mocified in certain cases., These criteria are application dependent and may need to be modified in certain cases.834 Generally speaking. our code uses two conditions: i) if quantities (like density or pressure)," Generally speaking, our code uses two conditions: i) if quantities (like density or pressure)"835With at accreting-ον Chemical evolution model aix some simiple assumptions. we find tLat τί is readily. possible to describe tle MDF of the Milkv Wav globular clusters. with au accuracy aud dLivsical basis superior o tlie staudard clouble-Gaussian uunerical fits that pervade the literaure.,"With an accreting-box chemical evolution model and some simple assumptions, we find that it is readily possible to describe the MDF of the Milky Way globular clusters, with an accuracy and physical basis superior to the standard double-Gaussian numerical fits that pervade the literature."836 An unavoidable. aud still rather arbitrary. key. [actor is that we are forced to adopt two clisinct dhases of cluster formation to prexluce the clear bimocality that the observations demauncl.," An unavoidable, and still rather arbitrary, key factor is that we are forced to adopt two distinct phases of cluster formation to produce the clear bimodality that the observations demand."837 Within his coutex. however. a wide range of model j»arameters can produce entirely adecuate its to the observations.," Within this context, however, a wide range of model parameters can produce entirely adequate fits to the observations."838 Tle conubination tha we [iud most »ersuasive is that: (a) The halo clusters fo1ued from jear-pΠοιοια [σας with a low vield yc20.00125 and shut down heir formation at an early stage (this 1wles the possibility that mauy of them formecl iu the stnall potential wells o- dwarf satelites and were accreted later). (, The combination that we find most persuasive is that: (a) The halo clusters formed from near-primordial gas with a low yield $y_{\rm eff} \simeq 0.0015$ and shut down their formation at an early stage (this includes the possibility that many of them formed in the small potential wells of dwarf satellites and were accreted later). (839b) The bulge clusers formed sta‘tine [rom uildly enriches eas (~0.22.) aud at yeyzz0.0OL.,b) The bulge clusters formed starting from mildly enriched gas $\sim 0.2 Z_{\odot}$ ) and at $y_{\rm eff} \simeq 0.0045$.840 Iu both cases. a siguilicaut phase of early eas infall is jecessary to reproduce the observed ΝDF shape.," In both cases, a significant phase of early gas infall is necessary to reproduce the observed MDF shape."841 To the extent tha two sharoly distinct phases rep'esent something real in the actual lisory ol the GCs (e.g.Santos2002).. his basi€ approach ueeds to ye explored. further.," To the extent that two sharply distinct phases represent something real in the actual history of the GCs \citep[e.g.][]{san02}, this basic approach needs to be explored further."842 Án imporat uew piece of the evolttionary plzzle uow emergiug is hat the MDFs of the globular clusters iu several galaxies are strikinely cdiflerent rol the metalicity «istributious of the halo aud bilge (Harris&Harris.200]L:Durrell.Pritche2001:Harris2002).," An important new piece of the evolutionary puzzle now emerging is that the MDFs of the globular clusters in several galaxies are strikingly different from the metallicity distributions of the halo and bulge \citep{hh01,dur01,har02}."843. lu ealaxies such as the LNC. 131. 132. ara4 NCC 25128 (aud perliaps in giant E galaxies generaly). the old-halo stellar populatiOW ls very ]road aud strongly weigued to moderately high metallicity near [Fe/H] ~—0.5.," In galaxies such as the LMC, M31, M32, and NGC 5128 (and perhaps in giant E galaxies generally), the old-halo stellar population is very broad and strongly weighted to moderately high metallicity near [Fe/H] $\sim -0.5$."844 TIese MDEs exhibit only a thin metal-poor tail aud uo trace of the disti1ict bimocality characterizit etle elobular clusters., These MDFs exhibit only a thin metal-poor tail and no trace of the distinct bimodality characterizing the globular clusters.845 This evidence reinforces the suspicion that the story oL formation for the eloular clusters back quite distinct elements., This evidence reinforces the suspicion that the story of formation for the globular clusters had quite distinct elements.846 Although we have discissec specific results ouly for the Milky Way. exteusions to other galaxies with clearly bimodal MDFs are obvious.," Although we have discussed specific results only for the Milky Way, extensions to other galaxies with clearly bimodal MDFs are obvious."847 This work was supported by the Natural Scieuces aud Eueineering Research Council of Canada through research grants to ΕΠ., This work was supported by the Natural Sciences and Engineering Research Council of Canada through research grants to WEH.848on the plivsical nature of the hieh-velocity feature.,on the physical nature of the high-velocity feature.849" llow. then. might one account for a principle polarization axis of SN 2001el. the dispersion about (hat axis and. especially, the detached. high-velocity feature seen so stronely in SN 2001el?"," How, then, might one account for a principle polarization axis of SN 2001el, the dispersion about that axis and, especially, the detached, high-velocity feature seen so strongly in SN 2001el?"850 Detailed modeling with 3-D radiative transfer is required. to fully. understand the polarization data (Ilóflichοἱal.1996:Wang.Wheeler.& 2001)..," Detailed modeling with 3-D radiative transfer is required to fully understand the polarization data \citep{Hoeflich:1996w, WWH:1997,851Howell:99by}."852 This analvsis will be presented elsewhere., This analysis will be presented elsewhere.853 Here we will discuss some of the physical possibilities., Here we will discuss some of the physical possibilities.854 As outlined in the Introduction. there are a variety of wavs that svstematie asvimnmetries could be imposed on the explosion of à SN la: rotation of the exploding star. an accretion disk. (he presence of a binary companion. or plumes of combustion products.," As outlined in the Introduction, there are a variety of ways that systematic asymmetries could be imposed on the explosion of a SN Ia: rotation of the exploding star, an accretion disk, the presence of a binary companion, or plumes of combustion products."855 Some of these could define the dominant. axis. others could account. for varving orientation axes and still others for the dispersion observed around (he dominant axis.," Some of these could define the dominant axis, others could account for varying orientation axes and still others for the dispersion observed around the dominant axis."856 Some affects of rotation have been sketched in Howell(2001) and we will not repeat them here.," Some affects of rotation have been sketched in \citet{Howell:2001}857 and we will not repeat them here."858 show that the impact of the supernova ejecta with the secondary star. assumed to fill its Roche lobe. creates a hole in the ejecta with an angular size of ~30 iin the high-velocity ejecta and with an angular size 40° in the low-velocity ejecta. or of the ejecta’s surface.," \citep{Marietta:2000} show that the impact of the supernova ejecta with the secondary star, assumed to fill its Roche lobe, creates a hole in the ejecta with an angular size of $\sim 30$ in the high-velocity ejecta and with an angular size $\sim40$ in the low-velocity ejecta, or of the ejecta's surface."859 This elfect might be able to induce the small photospheric polarization along the dominant axis (hat we report here. and it might not be observable in all SN Ia due to orientation effects.," This effect might be able to induce the small photospheric polarization along the dominant axis that we report here, and it might not be observable in all SN Ia due to orientation effects."860 We will return to the possible effects of an accretion disk below., We will return to the possible effects of an accretion disk below.861 The greatest challenge in the current observations is to account for the high-velocity shell ol ealeium-rich. material., The greatest challenge in the current observations is to account for the high-velocity shell of calcium-rich material.862 The low-velocitv Ca II seems to share the velocity. the degree of polarization. aud (he polarization angle with the photosphere.," The low-velocity Ca II seems to share the velocity, the degree of polarization, and the polarization angle with the photosphere."863 By contrast. the high-velocity malter defined by the 800 nm leature differs from the photosphere in velocity (by definition). in the degree of polarization. in the polarization angle. in the optical thickness. and in the filline factor.," By contrast, the high-velocity matter defined by the 800 nm feature differs from the photosphere in velocity (by definition), in the degree of polarization, in the polarization angle, in the optical thickness, and in the filling factor."864 IF the Ca II identification is correct. the high-velocity matter is physically and eeonmetricallv detached Irom the lower-velocily material on (he basis of the sharp edges of (he absorption lines.," If the Ca II identification is correct, the high-velocity matter is physically and geometrically detached from the lower-velocity material on the basis of the sharp edges of the absorption lines."865 The depth of hieh velocity feature requires that there must be (wo zones ol Ca II., The depth of high velocity feature requires that there must be two zones of Ca II.866 These zones differ bv so much in geometry. dvnamies and column density that it is difficult to accommodate them in an exploding star of only one component.," These zones differ by so much in geometry, dynamics and column density that it is difficult to accommodate them in an exploding star of only one component."867 The impressive homogeneitv of SN Ia does not leave much room for major individual peculiarities., The impressive homogeneity of SN Ia does not leave much room for major individual peculiarities.868 This is, This is869a beatiful confirmation of the results obtained from ground-based observations αἱ lower resolution.,a beautiful confirmation of the results obtained from ground-based observations at lower resolution.870 SainzDalda&BellotRubio(2008) then discovered that the opposite polarity patches move radially outwaid. often accompanied by another patch having the polarity of the spot and being located further away [rom the umbra.," \cite{dalda08} then discovered that the opposite polarity patches move radially outward, often accompanied by another patch having the polarity of the spot and being located further away from the umbra."871 SainzDalda&MartinezPillet(2005) and Ravindra(2006) also detected the evolution of reversed-polaritv patches in the penumbra using SOIIO/MDI observations (Scherreretal.1995).. but their bipolar nature could not be established.," \cite{dalda05} and \cite{ravindra06} also detected the evolution of reversed-polarity patches in the penumbra using SOHO/MDI observations \citep{scherrer1995}, but their bipolar nature could not be established."872 Taken together. these results provide an indication (hat the moving magnetic features observed in the penumbra and the Evershed flow belong to (he sume physical mechanism.," Taken together, these results provide an indication that the moving magnetic features observed in the penumbra and the Evershed flow belong to the same physical mechanism."873 The magnetic patches move between (he more vertical filaments of (he penumbra with a (vpical velocity of 0.3—1 km/s. their length is 23 aresec. (he mean width is ~1.5 aresec. and they have a lifetime of 0.5 7 hours (SainzDalda&BellotRubio2008).," The magnetic patches move between the more vertical filaments of the penumbra with a typical velocity of $0.3-1$ km/s, their length is 2–3 arcsec, the mean width is $\sim 1.5$ arcsec, and they have a lifetime of 0.5 – 7 hours \citep{dalda08}."874. Sometimes. the bipolar pairs appear one behind another along the same filament.," Sometimes, the bipolar pairs appear one behind another along the same filament."875 This can be (he signature of a wave-like behavior. similar to the wave behavior of magnetoconvection in (he penumbra found in numerical simulations (IXitiashvilietal.2009a.b).," This can be the signature of a wave-like behavior, similar to the wave behavior of magnetoconvection in the penumbra found in numerical simulations \citep{kiti09a,kiti09b}."876. Based on the observational data. SainzDalda&BellotRubio(2008) proposed a serpent model which explains the existence of moving bipolar süructures as wavering magnetic lied lines.," Based on the observational data, \citet{dalda08} proposed a sea-serpent model which explains the existence of moving bipolar structures as wavering magnetic field lines."877 In (his model. almost horizontal field lines return (o the solar interior and then come back to the photosphere. thus forming bipolar patches.," In this model, almost horizontal field lines return to the solar interior and then come back to the photosphere, thus forming bipolar patches."878 This scenario is similar to that discussed by. Schlichenmaier (2002) in the context of the moving tube model., This scenario is similar to that discussed by Schlichenmaier (2002) in the context of the moving tube model.879 It is also similar to the model proposed by Harvey&(1973). [or moving magnetic features in (he sunspot moat (MMESs:Sheeley1969)., It is also similar to the model proposed by \cite{harvey73} for moving magnetic features in the sunspot moat \citep[MMFs;][]{sheeley1969}.880. Indeed. it has been suggested that MMESs represent the continuation of similar features in (he penumbra (SainzDalda&MartinezRavindra2006:SainzDalda&BellotRubio2008:Kuboetal. 2005).," Indeed, it has been suggested that MMFs represent the continuation of similar features in the penumbra \citep{dalda05,ravindra06,dalda08,kubo08}. ."881. For numerical simulations of sunspot penumbra conditions we use (he 3D radiative MIID code “SolarBox” (Jacoutotοἱal.2008a:Jacoutotet2008b).," For numerical simulations of sunspot penumbra conditions we use the 3D radiative MHD code ""SolarBox"" \citep{jacoutot08a,jacoutot08b}."882. This code takes into account all essential physics and includes sub-erid scale turbulence modeling based on the lavee-eddy simulation (LES) approach., This code takes into account all essential physics and includes sub-grid scale turbulence modeling based on the large-eddy simulation (LES) approach.883 A dvnamie Smagorinsky turbulence model provides the best agreement wilh observations in terms of the acoustic oscillation power., A dynamic Smagorinsky turbulence model provides the best agreement with observations in terms of the acoustic oscillation power.884 However. in (hese simulaiions we use a computationally more efficient hyperviscosity model. because ib shows results qualitatively very similar to the dynamic model.," However, in these simulations we use a computationally more efficient hyperviscosity model, because it shows results qualitatively very similar to the dynamic model."885 The code uses a real-gas equation of state. Lakes into account ionization and excitation of all abundant species in the LTE approxination. audincludes radiative transfer and magnetic effects," The code uses a real-gas equation of state, takes into account ionization and excitation of all abundant species in the LTE approximation, andincludes radiative transfer and magnetic effects"886With a standard SMZ only. a deep as the SCZ. the time needed to erase the ssignature is of the order of 1 Myr (Tureotteand.Charbonneau1993).,"With a standard SMZ only a deep as the SCZ, the time needed to erase the signature is of the order of 1 Myr \citep{TC93}."887. The timescale in the case of deep mixing will be increased by a factor of some few hundreds because of the increase in (he mass of the SMZ. but (his is mitigated by the relative increase in [lux due to meridional circulation.," The timescale in the case of deep mixing will be increased by a factor of some few hundreds because of the increase in the mass of the SMZ, but this is mitigated by the relative increase in flux due to meridional circulation."888" The net effect for a mixed mass of LO° M, is an increase of the timescale [from by a [actor of 5 only.", The net effect for a mixed mass of $10^{-6}$ $_\star$ is an increase of the timescale from by a factor of 5 only.889 A complicating [actor [for the accretion scenario is (hat not only is it necessary (o dramatically increase (he amount of gas accreted on the star in order to impart (he composition of dust-depleted. circumstellar gas (ο the MZ. which may be accounted by much larger accretion rales on the pre-main-sequence. but much larger ongoing accretion rates are necessary to sustain the abundance peculiarities i£ (he 9MZ is as deep as suggested here.," A complicating factor for the accretion scenario is that not only is it necessary to dramatically increase the amount of gas accreted on the star in order to impart the composition of dust-depleted circumstellar gas to the SMZ, which may be accounted by much larger accretion rates on the pre-main-sequence, but much larger ongoing accretion rates are necessary to sustain the abundance peculiarities if the SMZ is as deep as suggested here."890 Fig. 2..," Fig. \ref{fig:mloss},"891" shows that an accretion rate of 10 to 10.1 NL,yr| is necessary to just. balance the flux of parlicles entering or leaving the SMZ at ils base from diffusion and meridional circulation."," shows that an accretion rate of $10^{-12}$ to $10^{-11}$ $_\sun\,yr^{-1}$ is necessary to just balance the flux of particles entering or leaving the SMZ at its base from diffusion and meridional circulation."892 such large rates may not problematic as they have been claimed in Pictoris 1996)., Such large rates may not problematic as they have been claimed in $\beta$ Pictoris \citep{Beustetal96}.893. It might however raise questions as to whether (he amount of circumstellar matter required (to provide such large rates could remain unseen as is the case in many sstars., It might however raise questions as to whether the amount of circumstellar matter required to provide such large rates could remain unseen as is the case in many stars.894" Still. if one assumes that it is the case and that the necessary accretion is ongoing as long as the circunistellar disk is present. one would still not expect sstars as old as 1 Gyr,"," Still, if one assumes that it is the case and that the necessary accretion is ongoing as long as the circumstellar disk is present, one would still not expect stars as old as 1 Gyr."895 We have restricted ourselves to rotational velocities of LOO ! or lower because of the limitations of the formalism for meridional cireulation used here., We have restricted ourselves to rotational velocities of 100 $^{-1}$ or lower because of the limitations of the formalism for meridional circulation used here.896 sstars can rotate at a much faster rate. as much as 250 | (Paunzen2001).," stars can rotate at a much faster rate, as much as 250 $^{-1}$ \citep{Paunzen01}."897. In such stars. the meridional cireulation would dominate a given accretion rate lor much shallower SMZs.," In such stars, the meridional circulation would dominate a given accretion rate for much shallower SMZs."898 The accretion rates required to establish the ssignature could then be an order of magnitude larger. or more. than those found for the models discussed here.," The accretion rates required to establish the signature could then be an order of magnitude larger, or more, than those found for the models discussed here."899" Finally. as sslars often are pulsating stars. it is tantalizing to imagine that (here might be a seismic signature of (he depth of the mixing considering that the abundance of most metals is completely different in (he ""metal opacity bump” depending on the models discussed here."," Finally, as stars often are pulsating stars, it is tantalizing to imagine that there might be a seismic signature of the depth of the mixing considering that the abundance of most metals is completely different in the “metal opacity bump” depending on the models discussed here."900" As the metals play a role in driving pulsations and determining the structure of the envelope in these stars. accretion with deep mixing might vield an observable signature in either which modes become overstable or in shifts in the Ireeuenucies of modes of pulsations will respect to standard models for sstars,"," As the metals play a role in driving pulsations and determining the structure of the envelope in these stars, accretion with deep mixing might yield an observable signature in either which modes become overstable or in shifts in the frequencies of modes of pulsations with respect to standard models for stars."901 Preliminary models have shown shifts in frequencies bv as much as 10 to but a seismic test for mixing in," Preliminary models have shown shifts in frequencies by as much as 10 to \citep{T00}902 but a seismic test for mixing in"903rest of the observation.,rest of the observation.904 The best-fit parameters of both spectra are consistent with each other., The best-fit parameters of both spectra are consistent with each other.905 Results of the X-ray spectral analysis can be found in Table I., Results of the X-ray spectral analysis can be found in Table 1.906 GRB 081029 is characterised by a complex optical and near-IR light-curve as shown in Fig. |.., GRB 081029 is characterised by a complex optical and near-IR light-curve as shown in Fig. \ref{lcopt}.907 We can divide the optical-NIR light-curve into three phases: In the following section we will describe the observed optical-NIR evolution during these three phases., We can divide the optical-NIR light-curve into three phases: In the following section we will describe the observed optical-NIR evolution during these three phases.908 The seven GROND optical-NIR light-curves of phase 1) can be well represented by a smoothly-connected broken power-law T=0.95) while a single power-law fit is excluded T-4.4)., The seven GROND optical-NIR light-curves of phase i) can be well represented by a smoothly-connected broken power-law $\chi^2_{\rm red}=0.95$ ) while a single power-law fit is excluded $\chi^2_{\rm red}=4.4$ ).909 If we contemporaneously fit these seven bands using the parametrisation form Beuermann et al. (, If we contemporaneously fit these seven bands using the parametrisation form Beuermann et al. (9101999) where T (a) is the pre-(post-)break decayindex?.. s' is the sharpness of the break. the apex (/) the light-curve phase we are describingp and the break timeη fιδ is- definedN as. We; obtain: αιu=0.38+0:05 and a?=1.12+0.06. and an achromatic break ín located at 940 + 30 s. The optical magnitudes observed by the REM telescope 2.5 min after the trigger (Covino et al.,"1999) where $\alpha_1^{(i)}$ $\alpha_2^{(i)}$ ) is the pre-(post-)break decay, $s^{(i)}$ is the sharpness of the break, the apex $(i)$ the light-curve phase we are describing and the break time ${t_{{\rm b}}^{(i)}}$ is defined as, We obtain $\alpha_1^{(i)}=0.38 \pm 0.05$ and $\alpha_2^{(i)}=1.12 \pm 0.06$, and an achromatic break ${t_{{\rm b}}^{(i)}}$ located at 940 $\pm$ 30 s. The optical magnitudes observed by the REM telescope 2.5 min after the trigger (Covino et al."911 2008a) are consistent. with the extrapolation to earlier times of this curve., 2008a) are consistent with the extrapolation to earlier times of this curve.912 À significant deviation from this model can be seen after 2.2 ks. immediately before the start of the intense rebrightening. where the light-curve commences a flattening.," A significant deviation from this model can be seen after 2.2 ks, immediately before the start of the intense rebrightening, where the light-curve commences a flattening."913 Unfortunately. we lack observations between 2.7 ks and 3.5 ks. exactly around the beginning of the rebrightening.," Unfortunately, we lack observations between 2.7 ks and 3.5 ks, exactly around the beginning of the rebrightening."914 A steep rise is observed in all seven GROND bands and starts between 2.5 ks and kks., A steep rise is observed in all seven GROND bands and starts between 2.5 ks and ks.915 The lack of observations during this interval does not allow us to precisely test the achromaticity of this start., The lack of observations during this interval does not allow us to precisely test the achromaticity of this start.916 Between 3.5 ks and 4.8 ks the light-curve brightens in all bands by more than 1.1 mag., Between 3.5 ks and 4.8 ks the light-curve brightens in all bands by more than 1.1 mag.917 This rise is very well tracked by the GROND photometry with twelve |-min observations in all seven bands., This rise is very well tracked by the GROND photometry with twelve 1-min observations in all seven bands.918 After a short constant flux state lasting about 400 s. around 5.2 ks after the trigger another rise of about 0.2 mag leads to the maximum at 5.9 ks (see Fig. 3)).," After a short constant flux state lasting about 400 s, around 5.2 ks after the trigger another rise of about 0.2 mag leads to the maximum at 5.9 ks (see Fig. \ref{modeltriple}) )."919 In this discussion we consider this further steep rebrightening at 5.2 ks as the first of a series of optical flares superposed on the post-break shallower power-law continuum., In this discussion we consider this further steep rebrightening at 5.2 ks as the first of a series of optical flares superposed on the post-break shallower power-law continuum.920 The brightness of the afterglow and the consequent small error bars of the optical photometry during the rise allow the identification of several substructures that make the light-curve deviate from a simple power-law., The brightness of the afterglow and the consequent small error bars of the optical photometry during the rise allow the identification of several substructures that make the light-curve deviate from a simple power-law.921" Considering only the data after ss. the steep rise requires a power-law index a2-4.7 and a break around ss. The position of the break is not well constrained because of the ""flare"" at kks."," Considering only the data after s, the steep rise requires a power-law index $\alpha_1^{(ii),(iii)}=-4.7$ and a break around s. The position of the break is not well constrained because of the “flare” at ks."922 If we take into account the possible contribution of the broken power-law, If we take into account the possible contribution of the broken power-law923covvariance.,variance.924comparison sample of solar-type stars spanr a wide range in stellar parameters (51OTN < < 6100 I. 3.6 < logy < L6.-1.6 < [Fe/II] < 10.2) and therefore hey are not representative of o1e-solar-mass solar anaoes.,"comparison sample of solar-type stars span a wide range in stellar parameters (5400 K $<$ $<$ 6100 K, 3.6 $<$ log $g$ $<$ 4.6, -1.6 $<$ [Fe/H] $<$ +0.2) and therefore they are not representative of one-solar-mass solar analogs."925 Iu Fig., In Fig.926 2. we restrict the comLISOLL siuuple of to only solar analogs within E2001. in solar effective temperature. +0.3 dex of the solar siface gravitv and 0.3 dex of he solar metallicity.," 2, we restrict the comparison sample of to only solar analogs within $\pm$ 200K in solar effective temperature, $\pm$ 0.3 dex of the solar surface gravity and $\pm$ 0.3 dex of the solar metallicity."927 As we cau see. these solar analogso seen to cluster in two groups. one wih very high lithiua abuudauces of Ap; ~ 2.4. i6. οι) times lig than solar. aud the other group with Li: vbunidances as low as solar.," As we can see, these solar analogs seem to cluster in two groups, one with very high lithium abundances of $A_{\rm Li}$ $\sim$ 2.4, i.e., 20 times higher than solar, and the other group with Li abundances as low as solar."928 Why are there no solar analoes with intermediate Li abundances?, Why are there no solar analogs with intermediate Li abundances?929 Why the uuaver of solar analoeso with low Li abundanuces is match lower than the umuber of analogs with high Li abuudawes’?, Why the number of solar analogs with low Li abundances is much lower than the number of analogs with high Li abundances?930 CotId this be he reason why oilv oud stars with hieh Li abuudanuces aLolnd «one solar lass?, Could this be the reason why only found stars with high Li abundances around one solar mass?931 The lack of stars with intermedia eL1 abundances in the Pasquini ct al. (, The lack of stars with intermediate Li abundances in the Pasquini et al. (932"199) solar analog salple. iux he lack of stars with both iiteriediae and low Li abuudaice around oue solar mass in the mde are probably telling us that bot1 snp lave biases, perhaps due ο a selection of stars mostly iu One or two evolutionary stages. e.g. mainly vouug stars. which are known to have high Li abundances.","1994) solar analog sample, and the lack of stars with both intermediate and low Li abundance around one solar mass in the sample are probably telling us that both samples have biases, perhaps due to a selection of stars mostly in one or two evolutionary stages, e.g., mainly young stars, which are known to have high Li abundances."933 Tje recent work by for solar analogs and solar twins iu the solar-age ope1 cluster M67. shows that solar twins (AIG7 stars around solar effecive temperature) have Li abunudaice as low as solar. but stars LOO or 200 Is) hotter span a broad range iu Li abundances.," The recent work by for solar analogs and solar twins in the solar-age open cluster M67, shows that solar twins (M67 stars around solar effective temperature) have Li abundance as low as solar, but stars 100 or 200 K hotter span a broad range in Li abundances."934 So. it is iuportant that the temperature scale of the compariso1 salple is accurate: otherwise offsets of about LOO I& may iutrocduce a bias in the comparison between the Sun aud stars.," So, it is important that the temperature scale of the comparison sample is accurate; otherwise offsets of about 100 K may introduce a bias in the comparison between the Sun and stars."935 Solar twins. stars witi stellar parameters very similar to the Sun. are ideal targets to see if the Sun is normal (or not) iu its Li abundance.," Solar twins, stars with stellar parameters very similar to the Sun, are ideal targets to see if the Sun is normal (or not) in its Li abundance."936 Beine so simular to the Suu. it is possible to obtain reliabe stellar parameters. and provided the Sun aud the twii are analyzed (aud observed) iu a consiseut wav. the temerature scale is accurate.," Being so similar to the Sun, it is possible to obtain reliable stellar parameters, and provided the Sun and the twins are analyzed (and observed) in a consistent way, the temperature scale is accurate."937 Furthermore. being selected. duc to their siuiluitv iu colors aud bIpmuuiuosiv to the Sun. they should span a range of ages very close to solar. avoiding thus potential biases iu the selectio1 of Comparison stars in only one evolutionary stage verv different to the preseut Sun.," Furthermore, being selected due to their similarity in colors and luminosity to the Sun, they should span a range of ages very close to solar, avoiding thus potential biases in the selection of comparison stars in only one evolutionary stage very different to the present Sun."938 Solar twins have been searched for a long time. aud although iuterestiug solar twin caudidates like 16 Ce D (ID 186127) were identi&ed iu the past. detailed analysis showed that they were sienificautlv different to the Sun1996).," Solar twins have been searched for a long time, and although interesting solar twin candidates like 16 Cyg B (HD 186427) were identified in the past, detailed analysis showed that they were significantly different to the Sun."939. When the first close solar twin (18 Sco) was found ).. at seemed to have a Li abuudance near solar. but much better data showed tlvt its Li abuudauce js actually hnree times higher tlvan solu.," When the first close solar twin (18 Sco) was found , it seemed to have a Li abundance near solar, but much better data showed that its Li abundance is actually three times higher than solar."940 One solar twin ds corainly not an accepable uumnber (for a coniparisonu νοποσα the Sun aud sars. πο a large survey of solar twijs was urgentlv neced.," One solar twin is certainly not an acceptable number for a comparison between the Sun and stars, so a large survey of solar twins was urgently needed."941 The two largest recent efforts for finding field solar twin stars are clue nudertakineg by the eroup of Y. Euseda and by our group2009)., The two largest recent efforts for finding field solar twin stars are being undertaking by the group of Y. Takeda and by our group.942. Buportaity. whenever possible. we are obtainius very high S/N or our sample stars. οςαπο otlicqwise onlv upper linLB.s cau be obtained for rei Li abundances.," Importantly, whenever possible, we are obtaining very high S/N for our sample stars, because otherwise only upper limits can be obtained for their Li abundances."943 Iudeed. as shown by iu their Fig.," Indeed, as shown by in their Fig."944 12. their data with S/N ~150 ca- ouly estiniate upper lanits for starswith Ap; « 1.5. ic.. rey can onlv reliably determi1C Li abuudauces whe- rev are three times higher than solu.," 12, their data with S/N $\sim$ 150 can only estimate upper limits for stars with $A_{\rm Li}$ $<$ 1.5, i.e., they can only reliably determine Li abundances when they are three times higher than solar."945 Qur solar “twin survey ld DOC1 performed mainly with the 2.71 telescope at MeDonald observatory in 1ο North and with the 6.51 Mlaecllan Clav. telescope at Las Campanas observaory in he South., Our solar twin survey has been performed mainly with the 2.7m telescope at McDonald observatory in the North and with the 6.5m Magellan Clay telescope at Las Campanas observatory in the South.946 We have also obtained some Week|WIRES dataiuthe North aud VLT|UVES aud TARPS data in the South., We have also obtained some Keck+HIRES data in the North and VLT+UVES and HARPS data in the South.947 Our data has been taken at R G0.000-11O00. and achieviug S/N = 200-1000.," Our data has been taken at R = 60,000-110,000 and achieving S/N = 200-1000."948 The first pilot data set taken at eck resulted in the discovery of the πόσοιid best solar twin. TID 985618. about a decade éter the discovery of the first solar twin Ls Sco.," The first pilot data set taken at Keck resulted in the discovery of the second best solar twin, HD 98618, about a decade after the discovery of the first solar twin 18 Sco."949 ΠΟ OSGLS secs to be a solar twin as good as 18 Sco. and. as this twin. it has also a Li aliudance three times higher than solar.," HD 98618 seems to be a solar twin as good as 18 Sco, and, as this twin, it has also a Li abundance three times higher than solar."950 Learning frou the experience 6: our pilot Ixeck observations. we improved our criteria to select the best solar twins. cinοσα adjusting our seale for an i»pareut zero-point problem2009).," Learning from the experience of our pilot Keck observations, we improved our criteria to select the best solar twins, empirically adjusting our scale for an apparent zero-point problem."951. This is probably the reason why our first solar twin runi at MeDonald was verv successful., This is probably the reason why our first solar twin run at McDonald was very successful.952 Besides coΠΙΟ) the solar twin nature of 18 Sco axl ΠΟ 98618. we icutified two additional solar twins. WIP 56918 and TWP 73815ΙΟ). both with a low Li abundance similar to solar.," Besides confirming the solar twin nature of 18 Sco and HD 98618, we identified two additional solar twins, HIP 56948 and HIP 73815, both with a low Li abundance similar to solar."953 IIIP569Is roniadnus to this date the star that most closely reseiubles the Sun. with a similar to solar within 1 Ix. as recently confirmed by usine Subaru|TDS observations.," HIP56948 remains to this date the star that most closely resembles the Sun, with a similar to solar within 10 K, as recently confirmed by using Subaru+HDS observations."954 The vear 2007 was very prolific for solar twin studies. besides the twins found by our eroup. reported the cISCOVCLY. of the fifth solar twin. WIP 110963.," The year 2007 was very prolific for solar twin studies, besides the twins found by our group, reported the discovery of the fifth solar twin, HIP 110963."955 Tutercstinely. this," Interestingly, this"956"fact. assuming Chat half of the barvons in the Universe are located in unvirialized filamentary WIIIM (Cen&Ostriker1999).. this mean cdensitv may be estimated as where QO),&0.05 is the barvonic fraction of the flat ACIDAL Universe. and * is the critical density (e.g..Fuk","fact, assuming that half of the baryons in the Universe are located in unvirialized filamentary WHIM \citep{cen99}, this mean density may be estimated as where $\Omega_{\rm b} \simeq 0.05$ is the baryonic fraction of the flat $\Lambda$ CDM Universe, and $\rho_{\rm cr} \simeq 10^{-29}$ $^{-3}$ is the critical density \citep[e.g.,][]{fuk04}."957ugita&Peebles2004).. The fact that implies that if the recurrent jet activity incase.. the source has to be located instead in a real void of the galaxy. and matter distribution. at the outskirts ol a filamentary WIIIM (e£.Subralimanyanetal.2008:Salouris2009).. since WIIIM is expected to constitute matter overdensities 10—30 times above the mean value (pi).," The fact that $\rho_{\textrm{\scriptsize ``inn''}}^{\rm J1420} \sim \langle \rho_{\rm b} \rangle$ implies that if the recurrent jet activity in, the source has to be located instead in a real void of the galaxy and matter distribution, at the outskirts of a filamentary WHIM \citep[cf.][]{sub08,saf09}, since WHIM is expected to constitute matter overdensities $10-30$ times above the mean value $\langle \rho_{\rm b} \rangle$."958 Exteuded lobes of GRGs. even though being in many cases relict structures. may be still expanding supersonically wilh respect to the ambient medium. due to (heir relatively hieh internal pressures aud low-densitv environments (Subralimanyanetal.2008).," Extended lobes of GRGs, even though being in many cases relict structures, may be still expanding supersonically with respect to the ambient medium, due to their relatively high internal pressures and low-density environments \citep{sub08}."959. With no σου constraints on the temperature of the non-virialized eas [ar away from groups and cluster of ealaxies. it is however hard (ο evaluate precisely the sound speed of IGM of interest.," With no good constraints on the temperature of the non-virialized gas far away from groups and cluster of galaxies, it is however hard to evaluate precisely the sound speed of IGM of interest."960" For the purpose of rough estimates we write where prox. Pray, and πουν=Lyκ105 IKIN are the pressure. density. aud temperature of (he gas surrounding the lobes."," For the purpose of rough estimates we write where $p_{\rm IGM}$, $\rho_{\rm IGM}$ and $T_{\rm IGM} \equiv T_6 \times 10^6$ K are the pressure, density, and temperature of the gas surrounding the lobes."961 Meanwhile. the sideway expansion of the lobes. driven by the cocoon's internal pressure. is roughly v;&(p./pisi)7. and the related Mach ΠΟΥ of a bow shock driven in the IGM by the expanding cocoon is My=ΌροςΟν.," Meanwhile, the sideway expansion of the lobes, driven by the cocoon's internal pressure, is roughly $v_{\ell} \simeq (p_{\rm c} /\rho_{\rm IGM})^{1/2}$, and the related Mach number of a bow shock driven in the IGM by the expanding cocoon is ${\cal M}_{\ell} = v_{\ell} / c_{\rm s,\,IGM}$."962" If the currently observed lobes of are due to a primary jel activity. ie. if these lobes evolve within ‘unelisturbed’ environment. (ie expected ambient medium (WIIIM) temperature is 0.1<7,SI (Cen&Ostriker 1999).."," If the currently observed lobes of are due to a primary jet activity, i.e. if these lobes evolve within `undisturbed' environment, the expected ambient medium (WHIM) temperature is $0.1 \lesssim T_6 \lesssim 1$ \citep{cen99}. ."963" This. together with pj;~(p) and pew(1—4x10P 7. as indicated by the modeling discussed in the previous seclions. gives us ος~0.010. and. My,~(20—120)."," This, together with $\rho_{\rm IGM} \sim \langle \rho_{\rm b} \rangle$ and $p_{\rm c} \sim (1-4) \times 10^{-14}$ $^{-2}$, as indicated by the modeling discussed in the previous sections, gives us $v_{\ell} \sim 0.01\,c$, and ${\cal M}_{\ell} \sim (20 - 120)$."964 If the observed lobes are due to the recurrent jel activity. instead. the involved shock Mach. number would be reduced due to the expected larger temperature of the IGA heated at the bow shock driven by the outer cocoon during the previous epoch ofthe jet activity.," If the observed lobes are due to the recurrent jet activity instead, the involved shock Mach number would be reduced due to the expected larger temperature of the IGM heated at the bow shock driven by the outer cocoon during the previous epoch ofthe jet activity."965 For this outer cocoon. taking again," For this outer cocoon, taking again"966Because we seek (ο identify the svstems of galaxies corresponding to the weak lensing peaks. we focus (he redshift survey on the red galaxy population.,"Because we seek to identify the systems of galaxies corresponding to the weak lensing peaks, we focus the redshift survey on the red galaxy population."967 These red objects preferentially populate clusters., These red objects preferentially populate clusters.968 This procedure also suppresses the blue foreground populations: (he weak lensing map is insensitive to svstenis al redshift less (han z0.05 (see Figure 9))., This procedure also suppresses the blue foreground populations; the weak lensing map is insensitive to systems at redshift less than $ z \simeq 0.05$ (see Figure \ref{fig:sensitivity.ps}) ).969 We acquired spectra for the objects with the IHLIectospec NAM el al., We acquired spectra for the objects with the Hectospec (Fabricant et al.970 1993. 2005). a 300-Liber robotic instrument mounted on the MMT from FebruaryMN 1. 2009 to April 27. 2009.," 1998, 2005), a 300-fiber robotic instrument mounted on the MMT from February 1, 2009 to April 27, 2009."971 The Hectospec observation planning software (Roll et al., The Hectospec observation planning software (Roll et al.972 1998) efficient acquisition of a pre-selected sample of galaxies., 1998) enables efficient acquisition of a pre-selected sample of galaxies.973 The software enables assignment of priorities as a function of galaxy. properties., The software enables assignment of priorities as a function of galaxy properties.974 The spectra cover the wavelength range 3500 10.000 wwith a resolution of 6À..," The spectra cover the wavelength range 3500 — 10,000 with a resolution of $\sim$ 6."975 Exposure times ranged n 0.75 — 1.5 hours., Exposure times ranged from 0.75 — 1.5 hours.976 We reduced the data with the standard IHectospec pipeline (Mink et al., We reduced the data with the standard Hectospec pipeline (Mink et al.977" m2007) and derived redshilts with RVSAO ναί Mink 1998) with templates constructed "" purpose (Fabricant et al.", 2007) and derived redshifts with RVSAO (Kurtz Mink 1998) with templates constructed for this purpose (Fabricant et al.978 2005)., 2005).979 Repeat observations vield robust estimates of the median error in ez where z is (he redshift and ο is the speed of light., Repeat observations yield robust estimates of the median error in $cz$ where $z$ is the redshift and $c$ is the speed of light.980 For emission line objects. (he median error (normalized by (L4+2))is27 kms I: the median for absorption line objects (again normalized by (1+:)) is 37 km |. (," For emission line objects, the median error (normalized by $(1 + z)$ ) is 27 km $^{-1}$; the median for absorption line objects (again normalized by $(1 + z$ )) is 37 km $^{-1}$. ("981Geller et al.,Geller et al.982 2010: Fabricant et al., 2010; Fabricant et al.983 2005)., 2005).984 There are 4541 galaxies in the region with a measured redshift: our Hectospec observations provide 4405 of these and the SDSS provides 136., There are 4541 galaxies in the region with a measured redshift; our Hectospec observations provide 4405 of these and the SDSS provides 136.985 In addition. we measured Lectospec redshifts for 54 of the 136 SDSS objects.," In addition, we measured Hectospec redshifts for 54 of the 136 SDSS objects."986 The SDSS objects are brighter and at lower redshift (han the others we observed with Hectospec., The SDSS objects are brighter and at lower redshift than the others we observed with Hectospec.987 Table 1 lists all of the objects and their redshilts., Table 1 lists all of the objects and their redshifts.988" The Table includes the SDSS ID (column 1). the J2000 right ascension (column 2). the J2000 declination (column 3). the SDSS Petrosian r4, magnitude (column 4). the SDSS liber color g—r (Column 5). the SDSS fiber color i—/ (column 6). the redshilt (column 7). the redshift error (column 3). and the redshift source (column 9)."," The Table includes the SDSS ID (column 1), the J2000 right ascension (column 2), the J2000 declination (column 3), the SDSS Petrosian $r_{petro}$ magnitude (column 4), the SDSS fiber color $g-r$ (column 5), the SDSS fiber color $r-i$ (column 6), the redshift (column 7), the redshift error (column 8), and the redshift source (column 9)."989 Our redshifts agree verv well with the 54 objects in common with SDSS: these objects range in redshift [rom 0.019 to 0.468., Our redshifts agree very well with the 54 objects in common with SDSS; these objects range in redshift from 0.019 to 0.468.990 At lower redshifts. the galaxies in common are emission line objects aad at higher redshift they are from the SDSS red galaxy sample.," At lower redshifts, the galaxies in common are emission line objects and at higher redshift they are from the SDSS red galaxy sample."991 The overlapping objects are brighter than the tvpical galaxy in our sample., The overlapping objects are brighter than the typical galaxy in our sample.992 The mean dillerence 0z/(1+2) — 34 kms !: the median difference is -0.8 kms !., The mean difference $\delta{z}/(1 + z)$ = 3.4 km $^{-1}$; the median difference is -0.8 km $^{-1}$ .993 The dispersion around the mean is 36.8 kis f., The dispersion around the mean is 36.8 km $^{-1}$.994 Our first D targets in the CITO2dee? field are galaxies with rpelro<21.3-- i>O4 and g—r>1.0.," Our first priority targets in the $^2$ field are galaxies with $r_{petro} \leq 21.3$, $r-i > 0.4$ and $g-r > 1.0$."995 dieThe bright limit is the SDSS completeness limit (rij; = 11. we observe only fainter for efficiency., The bright limit is the SDSS completeness limit $r_{petro}$ = 17.77); we observe only fainter galaxies for efficiency.996 Our target selection vields 3992 objects., Our target selection yields 3992 objects.997 In (his subsample. we obtained spectra for 3021 galaxies ancl 149 stars.," In this subsample, we obtained spectra for 3021 galaxies and 149 stars."998 In(he original selection. we include," Inthe original selection, we include"999observed in MWC 480. (he excess is also predicted to show strong ἐνρα emission features of CO and water. which are not observed.,"observed in MWC 480, the excess is also predicted to show strong $K$ -band emission features of CO and water, which are not observed."1000 The models futher predict prominent. CO and water features over a wide range in accretion rate (LO75—LOM.vr.|: Muzerolle et 22004: Calvet et 11991)., The models further predict prominent CO and water features over a wide range in accretion rate $10^{-8}-10^{-5}\Msunperyr$; Muzerolle et 2004; Calvet et 1991).1001 Emission is predicted at lower accretion rates and absorption at higher accretion rates., Emission is predicted at lower accretion rates and absorption at higher accretion rates.1002 The relative rarity of CO overtone and water features (in emission or absorption) from Herbig Ae stars suggests (hat something is missing from (he assumed physical-thermal-chemical structure of the disk in these models., The relative rarity of CO overtone and water features (in emission or absorption) from Herbig Ae stars suggests that something is missing from the assumed physical-thermal-chemical structure of the disk in these models.1003 Possibilities include non-LTE abundances. photodissociation. and missing sources of continuum opacity and ionization.," Possibilities include non-LTE abundances, photodissociation, and missing sources of continuum opacity and ionization."1004 If water emission does nol appear to be responsible for the excess detected inward οἱ the dust sublimation radius in MWC 480. what is the origin of the excess?," If water emission does not appear to be responsible for the excess detected inward of the dust sublimation radius in MWC 480, what is the origin of the excess?"1005 The lack of molecular emission might be explained by (he presence of a competing source of continuum opacity that produces the hot compact excess., The lack of molecular emission might be explained by the presence of a competing source of continuum opacity that produces the hot compact excess.1006 Possible addiGonal opacity sources are high temperature condensates. IL... and [ree-[ree emission.," Possible additional opacity sources are high temperature condensates, $^-$, and free-free emission."1007 Materials such as corundum (AlSQ4). hibonite (CaAlIj5O454). perovskite T1O4). and. eehlenite (CasAl2510;) are found in CI chondrites and have hieh sublimation temperatures of 16401300 Ix (Posch et 22007).," Materials such as corundum $_2$ $_3$ ), hibonite $_{12}$ $_{19}$ ), perovskite $_3$ ), and gehlenite $_2$ $_2$ $_7$ ) are found in CI chondrites and have high sublimation temperatures of 1640–1800 K (Posch et 2007)."1008 ∐∪∖∖⊽≼↲∖⇁≼↲↕⋅⋅⊔∐↲⋡∖↽∐≀↧↴∖↽≼↲↕↽≻∪∪↕⋅∫∖⋟−∣↽≻≀↧↴∐≼⇂≼↲∐∐⊳∖⊽⊳∖⊽↥∪∐≼↲∐∎↓≺∢↕≼↲∐≺∢↕≼↲⊳∖⊽↙↽⊋⋮⋔⋝∖ <]0? (Ilenning et1991?: also D. Sargent. personal communication). which makes it less likelv that thev will contribute significantly to the A-band opacity.," However, they have poor $K$ -band emission efficiencies $Q_{\rm abs} \lesssim 10^{-3}$ (Henning et; also B. Sargent, personal communication), which makes it less likely that they will contribute significantly to the $K$ -band opacity."1009 Graphite grains. often invoked to explain the bbunmp in the interstellar extinction curve. also have a high sublimation temperature 2000IXIX) at interstellar pressures (Ixrugel 2003: Salpeter 1977).," Graphite grains, often invoked to explain the bump in the interstellar extinction curve, also have a high sublimation temperature $> 2000$ K) at interstellar pressures (Krugel 2003; Salpeter 1977)."1010 However. graphite may be destroved by processes such as chenmüsputtering al much lower temperatures ~L200 KIN (Lenzuni. Gail. Henning 1995: Duschl. Gail. Tscharuuter 1996).," However, graphite may be destroyed by processes such as chemisputtering at much lower temperatures $\sim 1200$ K (Lenzuni, Gail, Henning 1995; Duschl, Gail, Tscharnuter 1996)."1011 So it is unclear whether it can explain the observed hot. compact excess in MWC 480.," So it is unclear whether it can explain the observed hot, compact excess in MWC 480."1012 Another possibility is that solids disappear through sublimation over a finite range in temperature (ancl disk radius)., Another possibility is that solids disappear through sublimation over a finite range in temperature (and disk radius).1013 For example. in considering the balance between (he solid and gas phases for silicates. Duschl et ((1996) found that (ihe gas and dust phases coexist over a range in temperature. wilh a fraction of silicates surviving (o temperatures ~2100 IxIx αἱ the densities of the inner disk region (~0.1 AAU).," For example, in considering the balance between the solid and gas phases for magnesium-iron silicates, Duschl et (1996) found that the gas and dust phases coexist over a range in temperature, with a fraction of silicates surviving to temperatures $\sim 2100$ K at the densities of the inner disk region $\sim 0.1$ AU)."1014 Aluminum-caleium-silicates may. be sustained to somewhat hieher temperatures., Aluminum-calcium-silicates may be sustained to somewhat higher temperatures.1015 Detailed modeling of this kind ma be needed io understand the conditions under which dust grains may contribute a modest. residual continuum opacity in the high temperature inner regions of disks., Detailed modeling of this kind may be needed to understand the conditions under which dust grains may contribute a modest residual continuum opacity in the high temperature inner regions of disks.1016In order to derive. knowledge of cosmology from. observations of the Large scale structure of the universe. we must accurately record the angular positions and clistances to billions of galaxies ancl then analyse their. statistical distribution.,"In order to derive knowledge of cosmology from observations of the large scale structure of the universe, we must accurately record the angular positions and distances to billions of galaxies and then analyse their statistical distribution."1017 “To determine distances in cosmology it is first necessary to have a fiducial cosmological model that allows the observer to calculate quantities such as the angular diameter clistance di(z2) and Hubble function {1(ς)., To determine distances in cosmology it is first necessary to have a fiducial cosmological model that allows the observer to calculate quantities such as the angular diameter distance $d_A(z)$ and Hubble function $H(z)$.1018 Inevitably the underlving. model the cosmologist has is incorrect in some wav., Inevitably the underlying model the cosmologist has is incorrect in some way.1019 This results in geometric distortions known as the Alceock-Paczvnski effect. (2)— -. distances measured along the line of sight. look dillerent. to. those measured. perpendicular to the line of sight (772).. ," This results in geometric distortions known as the Alcock-Paczynski effect \citep{1979Natur.281..358A} - distances measured along the line of sight look different to those measured perpendicular to the line of sight \citep{2005MNRAS.364..743N, 2011MNRAS.tmp.1599B, 2011arXiv1105.2037K}."1020The redshift of a galaxy. z. is not a true measure of a galaxw's distance but a measure of the recessional velocity. of the galaxy.," The redshift of a galaxy, $z$, is not a true measure of a galaxy's distance but a measure of the recessional velocity of the galaxy."1021 Galaxies. are involved in motions apart from the Llubble How. such as falling onto a cluster under gravity. these result in a distortion of the perceived distance to a galaxy (2)..," Galaxies are involved in motions apart from the Hubble flow, such as falling onto a cluster under gravity, these result in a distortion of the perceived distance to a galaxy \citep{1987MNRAS.227....1K}."1022 Because clusters grow. under eravity. through studying their growth at cilferent redshifts by isolating these redshift space distortions. we are able to learn how gravity behaves at dillerent epochs.," Because clusters grow under gravity, through studying their growth at different redshifts by isolating these redshift space distortions, we are able to learn how gravity behaves at different epochs."1023 This may help us identify deviations from Ceneral Relativity (???).. ," This may help us identify deviations from General Relativity \citep{2010PhRvD..81d3512S, 2009MNRAS.393..297P, 2005PhRvD..72d3529L}."1024It is well known that these two effects are degenerate in some regimes., It is well known that these two effects are degenerate in some regimes.1025 Dynamic anc geometric distortions can both have an enhancing effect on the galaxy. power spectrum in the line of sight., Dynamic and geometric distortions can both have an enhancing effect on the galaxy power spectrum in the line of sight.1026 1n order to identify and study these effects it is necessary to find features that are known to be the same physical size in the radial and. transverse directions. and the same size at all redshifts - a so called. cosmic ruler., In order to identify and study these effects it is necessary to find features that are known to be the same physical size in the radial and transverse directions and the same size at all redshifts - a so called cosmic ruler.1027 Unfortunately astrophysical objects are neither uniform nor large enough to perform such tests., Unfortunately astrophysical objects are neither uniform nor large enough to perform such tests.1028 The scale of the Barvon Acoustic Oscillations (D.XO). Le. the sound horizon at the epoch of last scattering. rs. (seen on the CMD sky and first detected in the SDSS LRG galaxy distribution by 2)) can be used as such a standard ruler to test our understanding of cosmology (for a recent review see ?)).," The scale of the Baryon Acoustic Oscillations (BAO), i.e. the sound horizon at the epoch of last scattering, $r_s$, (seen on the CMB sky and first detected in the SDSS LRG galaxy distribution by \cite{2005ApJ...633..560E}) ) can be used as such a standard ruler to test our understanding of cosmology (for a recent review see \cite{2009arXiv0910.5224B}) )."1029 This characteristic clustering scale. which emerges from physical processes in," This characteristic clustering scale, which emerges from physical processes in"1030changing from ~5% for SNe =5 to <2% for SNe >50.,changing from $\sim 5\%$ for SNe =5 to $< 2\%$ for SNe $>50$.1031 After such correction was applied. we computed again the average shear from the STEP] and STEP? simulations. and obtained a typical bias of ~ for SNe = 5.," After such correction was applied, we computed again the average shear from the STEP1 and STEP2 simulations, and obtained a typical bias of $\sim$ for SNe $=$ 5."1032 We then estimated the accuracy on the mass that can be obtained from an image with the same noise and depth as in the R-band SUBARU image., We then estimated the accuracy on the mass that can be obtained from an image with the same noise and depth as in the $R-$ band SUBARU image.1033 To this end we dropped the assumption on constant shear. and produced more realistic simulations: the effect on galaxy shapes by weak lensing from a galaxy cluster was produced using the code. that will be described in a separate paper (Huang et al..," To this end we dropped the assumption on constant shear, and produced more realistic simulations: the effect on galaxy shapes by weak lensing from a galaxy cluster was produced using the code, that will be described in a separate paper (Huang et al.,"1034 in preparation)., in preparation).1035 To summarize. the code takes as input a catalog of galaxies produced by the tool: it computes the shear produced by a standard mass profile (e.g. Navarro-Frenk-White. NFW hereafter) and applies it to the ellipticities of the galaxies behind the cluster.," To summarize, the code takes as input a catalog of galaxies produced by the tool; it computes the shear produced by a standard mass profile (e.g. Navarro-Frenk-White, NFW hereafter) and applies it to the ellipticities of the galaxies behind the cluster."1036 Such catalog is then used in the software. configured with the telescope parameters suitable for the SUBARU telescope and with the exposure time of the R-band image. producing a simulated image: the background rms of such image was set to be as close as possible to that of the real image.," Such catalog is then used in the software, configured with the telescope parameters suitable for the SUBARU telescope and with the exposure time of the $R-$ band image, producing a simulated image; the background rms of such image was set to be as close as possible to that of the real image."1037 We considered for the lens a range of masses at logήνΜο=13.5.14.14.5.15.0. and à NEW mass profile with c;=6.," We considered for the lens a range of masses at $\log M_{\rm vir}/M_\odot = 13.5, 14, 14.5, 15.0$, and a NFW mass profile with $c_{\rm vir} = 6$."1038 Each simulation was repeated 50 times for each mass value. randomly changing the morphology. position and redshift of the galaxies.," Each simulation was repeated 50 times for each mass value, randomly changing the morphology, position and redshift of the galaxies."1039 For each of these images. we run our lensing pipeline with the same configuration used for the real data.," For each of these images, we run our lensing pipeline with the same configuration used for the real data."1040 The density of background galaxies used for the lensing analysis was ~20 gals aremin., The density of background galaxies used for the lensing analysis was $\sim 20$ gals $^{-2}$ .1041 The fit of the mass was done as described in?) and ?:: the expressions for the radial dependence of tangential shear y; derived by ο. and ? were used. and the NFW parameters (My. Cir) Were derived using à maximum likelihood approach.," The fit of the mass was done as described in \citet{Radovich08} and \citet{Romano10}: : the expressions for the radial dependence of tangential shear $\gamma_{T}$ derived by \citet{Bartelmann96} and \citet{Wright00} were used, and the NFW parameters $M_{\rm vir}$, $c_{\rm vir}$ ) were derived using a maximum likelihood approach."1042 In addition. the 2D projected mass can be derived in a non-parametric way by aperture densitometry. where the mass profile of the cluster is computed by the Z statistics (2?):: The mass is estimated as Mj)(@))=πβιζ(θι Loi. and Gu is chosen so that (4s.Ayu)~0.," In addition, the 2D projected mass can be derived in a non-parametric way by aperture densitometry, where the mass profile of the cluster is computed by the $\zeta$ statistics \citep{fahlman, apj...497l..61c}: The mass is estimated as $M_{\rm ap}(\theta_1) = \pi \theta_1^2 \zeta (\theta_1) \Sigma_{\rm crit}$ , and $\theta_{\rm out}$ is chosen so that $\bar{\kappa} (\theta_2, \theta_{\rm out}) \sim 0$."1043 The average errors on mass estimate obtained m such way are displayed in Table 3.. showing that masses can be estimated within an uncertainty of < for M>10'M...," The average errors on mass estimate obtained in such way are displayed in Table \ref{tab:msimul}, showing that masses can be estimated within an uncertainty of $<$ for $M \ge 10^{14} M_\odot$."1044 Such accuracy only includes the contribute due to shape measurement and mass fitting method. but it does not include the uncertainty due to the selection of the lensed galaxies.," Such accuracy only includes the contribute due to shape measurement and mass fitting method, but it does not include the uncertainty due to the selection of the lensed galaxies."1045 Finally. the masses derived by aperture densitometry are ~ 1.3 higher than those obtained by mass fitting: this is in agreement with ?.. who find Mop/Msp=1.34 for virial overdensity.," Finally, the masses derived by aperture densitometry are $\sim$ 1.3 higher than those obtained by mass fitting: this is in agreement with \citet{Okabe10}, who find $M_{\rm 2D}/M_{\rm 3D}=1.34$ for virial overdensity."1046 One of the most critical source of systematic errors. which can lead to an underestimation of the true WL signal. is dilution of the distortion due to the contamination of the background galaxy catalog by unlensed foreground and cluster member galaxies (seee.g.?)..," One of the most critical source of systematic errors, which can lead to an underestimation of the true WL signal, is dilution of the distortion due to the contamination of the background galaxy catalog by unlensed foreground and cluster member galaxies \citep[see e.g.][]{Broadh05}."1047 The dilution effect increases as the cluster-centric distance decreases because the number density of cluster galaxies that contaminate the faint galaxy catalog ts expected to roughly follow the underlying density profile of the cluster., The dilution effect increases as the cluster-centric distance decreases because the number density of cluster galaxies that contaminate the faint galaxy catalog is expected to roughly follow the underlying density profile of the cluster.1048 Thus. correcting for the dilution effect is important to obtain unbiased. accurate constraints on the cluster parameters and mass profile.," Thus, correcting for the dilution effect is important to obtain unbiased, accurate constraints on the cluster parameters and mass profile."1049 As discussed by ???.. the selection of background galaxies to be used for the weak lensing analysis can be done taking those galaxies redder than the cluster red sequence.," As discussed by \citet{Broadh05,Okabe10,Oguri10}, the selection of background galaxies to be used for the weak lensing analysis can be done taking those galaxies redder than the cluster red sequence."1050 However. such selection produces a low number density (10 galaxies/aremin- in our case). and correspondingly high uncertainties in the derived. parameters.," However, such selection produces a low number density (10 $^2$ in our case), and correspondingly high uncertainties in the derived parameters."1051 In the following. we compare the results obtained by different methods.," In the following, we compare the results obtained by different methods."1052 We first assumed that no information on the redshift is available. and that photometry from only one band is available (magnitude cut). or from more than two bands (color selection).," We first assumed that no information on the redshift is available, and that photometry from only one band is available (magnitude cut), or from more than two bands (color selection)."1053 Finally. we included in our analysis the photometric redshifts.," Finally, we included in our analysis the photometric redshifts."1054 The density of background galaxies is ~ 25-30 galaxies/arcmin?. see Table 4..," The density of background galaxies is $\sim$ 25-30 $^2$, see Table \ref{tab:nfw}."1055" In order to derive the mass. we need to know the critical surface density: D,,.. D,. and D, being the angular distances between lens and source. observer and source. and observer and lens respectively."," In order to derive the mass, we need to know the critical surface density: $D_{ls}$, $D_s$, and $D_l$ being the angular distances between lens and source, observer and source, and observer and lens respectively."1056 This quantity should be computed for each lensed galaxy., This quantity should be computed for each lensed galaxy.1057 As the reliability of photometric redshifts for the faint background galaxies is not well known. we prefer to adopt the single sheet approximation. where all background galaxies are assumed to lie at the same redshift. defined as (BC).," As the reliability of photometric redshifts for the faint background galaxies is not well known, we prefer to adopt the single sheet approximation, where all background galaxies are assumed to lie at the same redshift, defined as $\beta(z_s) = \left\langle \beta(z) \right\rangle $ ."1058 In the case of the selection based only on magnitude or colors. such value was derived from the COSMOS catalog of photometric redshifts (2).. to which the same cuts used for the Abell 383 catalog are applied.," In the case of the selection based only on magnitude or colors, such value was derived from the COSMOS catalog of photometric redshifts \citep{capak07}, to which the same cuts used for the Abell 383 catalog are applied."1059 Later on. we computed 5(z.) from the photometric redshifts themselves. and compared such two values.," Later on, we computed $\beta(z_s)$ from the photometric redshifts themselves, and compared such two values."1060 The mass was computed both by fitting a NFWprofile (Map=M). and byaperture densitometry (Mp).," The mass was computed both by fitting a NFWprofile $M_{\rm 3D} = M_{\rm vir}$), and byaperture densitometry $M_{\rm 2D}$ )."1061 In the case, In the case1062especially at long wavelength. are affected by the whole Sun. errors in the inferred spatial structure are strongly correlated.,"especially at long wavelength, are affected by the whole Sun, errors in the inferred spatial structure are strongly correlated."1063 An inversion for |N(r). like the inversion for etr). will have uncertainties in both NN and (he position r.," An inversion for $N(r)$, like the inversion for $c_{\rm ad}(r)$, will have uncertainties in both $N$ and the position $r$."1064 The derived profiles (11) will have errors compounded from all these sources: ο) INGr). and r. and thus larger than errors in any one.," The derived profiles (11) will have errors compounded from all these sources: $c_{\rm ad}(r)$, $N(r)$, and $r$, and thus larger than errors in any one."1065" As the SSA evolves. V [alls while V,, increases from zero. the core becoming more concentrated."," As the SSM evolves, $\nabla$ falls while $\nabla_\mu$ increases from zero, the core becoming more concentrated."1066 Helioseismology is adiabatie and non-evolutionary and cannot give the thermal or chemical structure without further assumptions or a full S$M., Helioseismology is adiabatic and non-evolutionary and cannot give the thermal or chemical structure without further assumptions or a full SSM.1067 Many 59M core features can be understood via mechanical-thermal core homology. without a detailed model (DlIudman Kennedy 1996)., Many SSM core features can be understood via mechanical-thermal core homology without a detailed model (Bludman Kennedy 1996).1068 Homology applied to theonly is valid for power-law opacity and Iuminositv generation and does require a polvtrope., Homology applied to the is valid for power-law opacity and luminosity generation and does require a polytrope.1069 A polvtrope is not even approximately valid for the present 55M. core in anv case (kennedy DBludinan 1999)., A polytrope is not even approximately valid for the present SSM core in any case (Kennedy Bludman 1999).1070 The homology in reality is violated somewhat. as (he exponents are nol constant aud the huminositv is not exclusively produced by one reaction But an accurate thermal and chemical reconstruction is not much less complicated ancl requires no fewer assumptions than a full 55M.," The homology in reality is violated somewhat, as the exponents are not constant and the luminosity is not exclusively produced by one reaction But an accurate thermal and chemical reconstruction is not much less complicated and requires no fewer assumptions than a full SSM."1071 On the other hand. the mechanical structure of a ΑΓ can be specified by the reconstruction of Section 2. rather than calculated.," On the other hand, the mechanical structure of a SSM can be specified by the model-independent reconstruction of Section 2, rather than calculated."1072 As p and P are in principle inferable from helioseismology. the numerical precision of the reconstruction can be estimated.," As $\rho$ and $P$ are in principle inferable from helioseismology, the numerical precision of the reconstruction can be estimated."1073 No ganode measurements are now available. so the 55M houvaney [requencey is used here.," No $g$ -mode measurements are now available, so the SSM bouyancy frequency is used here."1074" Combining ο). dead(r)/dr. and the SSAT ΝΕ} and assuming Py, = 5/3. the profile D(r) is obtained (9) and hence DC)."," Combining $c_{\rm ad}(r)$, $dc_{\rm ad}(r)/dr$, and the SSM $N(r)$ and assuming $\Gamma_{\rm ad}$ = 5/3, the profile ${\cal D}(r)$ is obtained (9) and hence $\Gamma (r)$."1075 Figure 1 shows the measured sound speed below the convection zone. the points being fixed bv helioseismic inversion. (Christensen-Dalsgaard 1997).," Figure 1 shows the measured sound speed below the convection zone, the points being fixed by helioseismic inversion (Christensen-Dalsgaard 1997)."1076 The BP9S $S5M D and Α profiles. respectively. ave shown in Figs.," The BP98 SSM $\Gamma$ and $N$ profiles, respectively, are shown in Figs."1077 2(a.b).," 2(a,b)."1078 In. Figure 3 is displaved (he ratio of the reconstructed to the 59M V(r)., In Figure 3 is displayed the ratio of the reconstructed to the SSM $\Gamma (r)$.1079 The helioseismic inversion points are separated by steps of A(r/R.)x 0.01. and we should thus expect agreement with the SSAL at level of about one percent. consistent with Fie.," The helioseismic inversion points are separated by steps of $\Delta(r/R_\odot )\simeq$ 0.01, and we should thus expect agreement with the SSM at level of about one percent, consistent with Fig."1080 3., 3.1081" The inner boundary al ο, = 0.05 marks (he terminus of (he present helioseismic inversion. (he outer boundary (he base of (he convective zone al Προ... = 0.71."," The inner boundary at $r/R_\odot$ = 0.05 marks the terminus of the present helioseismic inversion, the outer boundary the base of the convective zone at $r_{\rm BCZ}/R_\odot$ = 0.71."1082 Cutting olf the fit at both ends introduces artificial errors. ancl (hus the first three and last three points are not shown in Fig.," Cutting off the fit at both ends introduces artificial errors, and thus the first three and last three points are not shown in Fig."1083 3., 3.1084for gamma-ray burst sources (Paezviisski 1991. 2001: Haensel et al.,"for gamma-ray burst sources (Paczyńsski 1991, 2001; Haensel et al."1085 1991)., 1991).1086 ‘This is reason enough to study such coalescences., This is reason enough to study such coalescences.1087 However. in this numerical stucly we address only one specific question: how much quark matter. and with what velocities. is cjected when a quark star coalesces with a black hole.," However, in this numerical study we address only one specific question: how much quark matter, and with what velocities, is ejected when a quark star coalesces with a black hole."1088 The interest here is in the speculations that such ejecta may convert all neutron stars to quark stars., The interest here is in the speculations that such ejecta may convert all neutron stars to quark stars.1089 3odmer (1971) and Witten (1984) have conjectured that a three-Davor quark [uid is the ground state of hadronie matter., Bodmer (1971) and Witten (1984) have conjectured that a three-flavor quark fluid is the ground state of hadronic matter.1090 This (Cup. down and) strange quark matter. if cold. would be stable in the bulk at zero pressure fission into individual hadrons or (hyper) nuclei would only occur for the tiniest. specks of quark matter. composed. of less than a few thousand quarks each. (Farhi and Jaffe. 1984).," This (up, down and) strange quark matter, if cold, would be stable in the bulk at zero pressure---spontaneous fission into individual hadrons or (hyper) nuclei would only occur for the tiniest specks of quark matter, composed of less than a few thousand quarks each (Farhi and Jaffe, 1984)."1091 As Witten (1984) was quick to point out. a sullicientIy large (sell-gravitating) quantity of quark matter would be remarkably similar to conventional neutron starsa solar mass quark star would be about 10 km across. and the maximum mass of a quark star stable against collapse to a black hole would be about 2M..," As Witten (1984) was quick to point out, a sufficiently large (self-gravitating) quantity of quark matter would be remarkably similar to conventional neutron stars—a solar mass quark star would be about 10 km across, and the maximum mass of a quark star stable against collapse to a black hole would be about $2M_\odot$."1092" The TOV equation for quark stars of masses up to the maximum value has also been solved by Itoh. (1970): Brecher ancl Caporaso (1976): llaensel. Zelunik ancl Schaelfer. (1986): as well as by Aleock. Farhi ane Olinto (1986a) who also give a detailed. discussion of the properties of these ""strange"" stars and of the astrophysical context."," The TOV equation for quark stars of masses up to the maximum value has also been solved by Itoh (1970); Brecher and Caporaso (1976); Haensel, Zdunik and Schaeffer (1986); as well as by Alcock, Farhi and Olinto (1986a), who also give a detailed discussion of the properties of these “strange” stars and of the astrophysical context."1093 For recent reviews see Cheng. Dai and Lu (1998) ancl Madsen (1999).," For recent reviews see Cheng, Dai and Lu (1998) and Madsen (1999)."1094 The existence of self-bound quark matter and of quark stars remains a hypothesis., The existence of self-bound quark matter and of quark stars remains a hypothesis.1095 Even so. we are now winessing a revival in their theoretica study. prompted no doubt by acvaric‘es in computational techniques anc in the maturing of X-ray and gamma-ray astronomy. as well as by a deeper understanding of collective effects in cuark matter (ee. Alford. Rajagopa and Wilezek 1998: Rapp et al.," Even so, we are now witnessing a revival in their theoretical study, prompted no doubt by advances in computational techniques and in the maturing of X-ray and gamma-ray astronomy, as well as by a deeper understanding of collective effects in quark matter (e.g., Alford, Rajagopal and Wilczek 1998; Rapp et al."1096 1998)., 1998).1097 Several groups have computed the structure of rotating quark stars in full general relativity. and ciscussed their external metric in the context of kllz quasi-periodic oscillations (QPOs) observed in certain X-ray binaries (Gourgoulbon ct al.," Several groups have computed the structure of rotating quark stars in full general relativity, and discussed their external metric in the context of kHz quasi-periodic oscillations (QPOs) observed in certain X-ray binaries (Gourgoulhon et al."1098 1999. Stergioulas et al.," 1999, Stergioulas et al."1099 1999. Gondek-Itosisska et al.," 1999, Gondek-Rosińsska et al."1100 2001. Bombaci et a.," 2001, Bombaci et al."1101 2000). and an even greater number have investigate the possible connection between quark stars and cnerectic phenomena such as gamma-ray bursts (Paczvisski LOOL: Hlaensel. Paczviisski ane Amsterdamski 1991: Cheng and Dai 1996). soft gamma-repeaters. aka.," 2000), and an even greater number have investigated the possible connection between quark stars and energetic phenomena such as gamma-ray bursts (Paczyńsski 1991; Haensel, Paczyńsski and Amsterdamski 1991; Cheng and Dai 1996), soft gamma-repeaters, a.k.a."1102 SGRs (Alcock. Farhi and Olinto 1986b: Llorvath et al.," SGRs (Alcock, Farhi and Olinto 1986b; Horvath et al."1103 1993: Cheng anc Dai 19958. 2002: Dai and Lu 1998: Zhang. Xu and Qiao 2000: Usov 2001). and radio pulsars (Xu et al.," 1993; Cheng and Dai 1998, 2002; Dai and Lu 1998; Zhang, Xu and Qiao 2000; Usov 2001), and radio pulsars (Xu et al."1104 1999)., 1999).1105 his reference list is far (rom exhaustive., This reference list is far from exhaustive.1106 llowever. there is a shadow over all this activity.," However, there is a shadow over all this activity."1107" The Galaxy must not be contaminated with ""seeds? of quark matter.", The Galaxy must not be contaminated with “seeds” of quark matter.1108 Such seeds present no danger to the Earth. because ions are repelled by the high (~10 MeV) Coulomb barrier. surrounding quark matter of the type discussed. here (Farhi and Jalle 1984: Alcock et al.," Such seeds present no danger to the Earth, because ions are repelled by the high $\sim 10\,$ MeV) Coulomb barrier surrounding quark matter of the type discussed here (Farhi and Jaffe 1984; Alcock et al."1109 1986a), 1986a).1110 Jul neutrons are easily absorbed, But neutrons are easily absorbed1111Some short GRBs show gamma-ray emission with photon energies above 100 MeV (Abdoetal..2009) but it has to be noted that only few GRBs are detected in this energy band (e.g.Zhangetal..2011).,Some short GRBs show gamma-ray emission with photon energies above 100 MeV \citep{abdo2009} but it has to be noted that only few GRBs are detected in this energy band \citep[e.g.][]{zhang2011}.1112. For the case of GRBO090510 the high energy component had an isotropic luminosity of about 3x107 erg (Abdoetal..2009)., For the case of GRB090510 the high energy component had an isotropic luminosity of about $4 \times 10^{52}$ erg \citep{abdo2009}.1113. This high-energy luminosity can result in a comparable fluence of high-energy gamma-ray photons as has been estimated for the cosmic ray fluence in the previous paragraph., This high-energy luminosity can result in a comparable fluence of high-energy gamma-ray photons as has been estimated for the cosmic ray fluence in the previous paragraph.1114 Consequently. also the energy deposition of high energy gamma-ray photons can equal the galactic cosmic ray energy deposition integrated over about 100 years 1f such a burst ts located at a distance of | kpe.," Consequently, also the energy deposition of high energy gamma-ray photons can equal the galactic cosmic ray energy deposition integrated over about 100 years if such a burst is located at a distance of 1 kpc."1115 However. again its quite unlikely that high-energy photons will leave measurable traces in the geological record in form of radioactive nuclei.," However, again its quite unlikely that high-energy photons will leave measurable traces in the geological record in form of radioactive nuclei."1116 To conclude. even for advantageous estimates for the GRB rate in globular cluster its quite unlikely that these GRBs have left anomalies of radioactive isotopes in the geological record.," To conclude, even for advantageous estimates for the GRB rate in globular cluster its quite unlikely that these GRBs have left anomalies of radioactive isotopes in the geological record."1117 However. for a short period of time (<100 years) they could have elevated the level of radioactivity which. could. cause biological mutations leading to fast appearance of new species.," However, for a short period of time $\lesssim 100$ years) they could have elevated the level of radioactivity which could cause biological mutations leading to fast appearance of new species."1118 Nearby GRBs launched in globular clusters offer the exciting possibility to calculate back the approximate time of their occurrence., Nearby GRBs launched in globular clusters offer the exciting possibility to calculate back the approximate time of their occurrence.1119 This can be done due to the fact that globular clusters follow quite well defined orbits around the galactic center (Allenetal..2006) and therefore the time of encounters between the solar system and a specific globular cluster can be determined (seee.g.VandePutte&Cropper.2009)., This can be done due to the fact that globular clusters follow quite well defined orbits around the galactic center \citep{allen2006} and therefore the time of encounters between the solar system and a specific globular cluster can be determined \citep[see e.g.][]{vandeputte2009}.1120. Here a globular cluster encounter is defined as a time period when the distance between the solar system and the globular cluster is smaller than di., Here a globular cluster encounter is defined as a time period when the distance between the solar system and the globular cluster is smaller than $d_\mathrm{min}$.1121 With the knowledge of the fraction of time where at least one globular cluster can be found within dyin (see Sec. 2)), With the knowledge of the fraction of time where at least one globular cluster can be found within $d_\mathrm{min}$ (see Sec. \ref{sec:mindist}) )1122 and with adopting the duration of a typical globular cluster — Earth encounter the number of such encounters can be estimated for a specific time interval., and with adopting the duration of a typical globular cluster – Earth encounter the number of such encounters can be estimated for a specific time interval.1123 In Sec., In Sec.1124" 2. it was estimated that during a fraction of about 0.1(RGc/30Gpeγι) of the time considered. one globular cluster is found within the minimal distance dj, for a GRB occurrence."," \ref{sec:mindist} it was estimated that during a fraction of about $0.1 \, (R_{GC}/30\, \mathrm{Gpc}^{-3}\mathrm{yr}^{-1})$ of the time considered, one globular cluster is found within the minimal distance $d_\mathrm{min}$ for a GRB occurrence."1125 A typical encounter will last for O(107) years if an encounter length of «1 kpe and a relative velocity between the earth and the globular cluster of ~100 kmss! is assumed., A typical encounter will last for $\mathcal{O}(10^7)$ years if an encounter length of $\sim$ 1 kpc and a relative velocity between the earth and the globular cluster of $\sim$ 100 $^{-1}$ is assumed.1126 If now a GRB rate Roc=Gpe?yi! is considered this. results in a total time of 10* years of presence of at least one globular cluster within di; in the last Gyr.," If now a GRB rate $R_{GC} = 30\, \mathrm{Gpc}^{-3}\mathrm{yr}^{-1}$ is considered this results in a total time of $10^8$ years of presence of at least one globular cluster within $d_\mathrm{min}$ in the last Gyr."1127 With a typical encounter duration of ~107 years O(10) globular cluster — Earth encounters within a distance of diy are expected., With a typical encounter duration of $\sim10^7$ years $\mathcal{O}(10)$ globular cluster – Earth encounters within a distance of $d_\mathrm{min}$ are expected.1128 The expected probability for à GRB occurring in a globular cluster during an encounter can be estimated from the rate of GRBs per globular cluster and the duration of an globular cluster — Earth encounter., The expected probability for a GRB occurring in a globular cluster during an encounter can be estimated from the rate of GRBs per globular cluster and the duration of an globular cluster – Earth encounter.1129 If for this rate a value of @D=(Rec[30Gpe“yr!) per globular cluster per year (see Sec. 2))," If for this rate a value of $\Phi \approx 10^{-8}\,(R_{GC}/30\, 1130\mathrm{Gpc}^{-3}\mathrm{yr}^{-1})$ per globular cluster per year (see Sec. \ref{sec:mindist}) )"1131 is adopted and a typical duration of a globular cluster passage of 10 years is assumed. then the resulting probability is Dx107years=O.1(Rec/30Gpe“yr7!).," is adopted and a typical duration of a globular cluster passage of $10^7$ years is assumed, then the resulting probability is $\Phi \times 10^7\, \mathrm{years} \approx 0.1\,(R_{GC}/30\, 1132\mathrm{Gpc}^{-3}\mathrm{yr}^{-1})$."1133 For a GRB rate of Reo=30Gpcyr! on average one GRB is expected every 10 globular cluster — Earth encounters.," For a GRB rate of $R_{GC} = 30\, \mathrm{Gpc}^{-3}\mathrm{yr}^{-1}$ on average one GRB is expected every 10 globular cluster – Earth encounters."1134 Thus for the time span of the geological record one GRB ts likely during a globular cluster - solar system encounter., Thus for the time span of the geological record one GRB is likely during a globular cluster - solar system encounter.1135 In the above considerations it is assumed that the GRB rate is the same for each globular cluster., In the above considerations it is assumed that the GRB rate is the same for each globular cluster.1136 However. some globular clusters are especially prolific. producers of close binaries (Pooley&Hut.2006).," However, some globular clusters are especially prolific producers of close binaries \citep{pooley2006}."1137. These objects likely exhibit also a higher GRB rate (Grindlayetal..2006:Lee2010).," These objects likely exhibit also a higher GRB rate \citep{grindlay2006,lee2010}."1138 Consequently. the passage of a globular cluster hosting a large number of stellar binaries close to Earth can be identified as a potential time interval of a nearby GRB explosion.," Consequently, the passage of a globular cluster hosting a large number of stellar binaries close to Earth can be identified as a potential time interval of a nearby GRB explosion."1139 As a result of the above considerations. specific time intervals corresponding to globular cluster encounters can be identified for which the geological record can be searched for potential terrestrial signatures connected to a nearby GRB.," As a result of the above considerations, specific time intervals corresponding to globular cluster encounters can be identified for which the geological record can be searched for potential terrestrial signatures connected to a nearby GRB."1140 Furthermore correlations between events of mass extinction in the history of life (Raup&Sepkoski.1982) or periods of rapid development of new species (e.g.CambrianexplosionMarshall.2006) with periods of globular cluster encounters can be explored., Furthermore correlations between events of mass extinction in the history of life \citep{raup1982} or periods of rapid development of new species \citep[e.g. Cambrian explosion][]{marshall2006} with periods of globular cluster encounters can be explored.1141" With a few globular cluster passages during the last 600 Myr and a typical duration of O10"") years for each passage a fair fraction of chance coincidence between events of mass extinctions and globular cluster passages can be expected.", With a few globular cluster passages during the last 600 Myr and a typical duration of $\mathcal{O}(10^7)$ years for each passage a fair fraction of chance coincidence between events of mass extinctions and globular cluster passages can be expected.1142 Considering only passages of globular clusters hosting an elevated number of close binaries can help to reduce the probability of such a chance coincidence., Considering only passages of globular clusters hosting an elevated number of close binaries can help to reduce the probability of such a chance coincidence.1143 In Sec., In Sec.1144 4 it has been pointed out that periods with globular cluster passages close to the solar systems can be determined., \ref{sec:asso} it has been pointed out that periods with globular cluster passages close to the solar systems can be determined.1145 During these time periods nearby GRB events launched in globular clusters could have happened., During these time periods nearby GRB events launched in globular clusters could have happened.1146 For determining the impact of nearby GRBs on the biota by correlating globular cluster passages and events of mass extinction. it Is necessary to reliably calculate the globular cluster distance for the last 600 Myrs.," For determining the impact of nearby GRBs on the biota by correlating globular cluster passages and events of mass extinction, it is necessary to reliably calculate the globular cluster distance for the last 600 Myrs."1147 An accurate determination of globular cluster positions appear in principle possible since the time span of 600 Myrs corresponds to only a few globular cluster orbits around the galactic center that last typically O(100) Myrs., An accurate determination of globular cluster positions appear in principle possible since the time span of 600 Myrs corresponds to only a few globular cluster orbits around the galactic center that last typically $\mathcal{O}(100)$ Myrs.1148 Assuming a velocity of the globular clusters of O(100) kmss! an accuracy of about for the 3d motion would be needed to constrain their position to ~1 kpe for the last Gyr., Assuming a velocity of the globular clusters of $\mathcal{O}$ (100) $^{-1}$ an accuracy of about for the 3d motion would be needed to constrain their position to $\sim1$ kpc for the last Gyr.1149 However. with the current precision. reliable calculations of the cluster position with respect to the Earth can only be calculated back for about 50 Myrs ago (seeVandePutte&Cropper. 2009)..," However, with the current precision reliable calculations of the cluster position with respect to the Earth can only be calculated back for about 50 Myrs ago \citep[see][]{vandeputte2009}."1150 The required accuracy for the 3d motion of galactic globular clusters will likely be available in the foreseen future., The required accuracy for the 3d motion of galactic globular clusters will likely be available in the foreseen future.1151 The up-coming GAIA will provide unprecedented positional and radial velocity measurement of about one billion stars in the Milky Way., The up-coming GAIA will provide unprecedented positional and radial velocity measurement of about one billion stars in the Milky Way.1152 With these data it will be possible to, With these data it will be possible to1153The target clusters were observed with the multi-object instrument FLAMES (Pasquini et al. 2000) ,The target clusters were observed with the multi-object instrument FLAMES (Pasquini et al. \cite{P00}) )1154on VLT/UT2 (ESO. Chile).," on VLT/UT2 (ESO, Chile)."1155 The fiber link to UVES was used to obtain high-resolution spectra (R=45.000) for RGB and red clump objects.," The fiber link to UVES was used to obtain high-resolution spectra $R=45,000$ ) for RGB and red clump objects."1156 For all the clusters we performed observations with the CD3 cross-disperser. covering the wavelength range ~4750- 6800 Az: for Cr 261 and Mel 66 we also obtained spectra using the CD4 grating (~ 6600-10600 A)).," For all the clusters we performed observations with the CD3 cross-disperser, covering the wavelength range $\sim$ 4750--6800 ; for Cr 261 and Mel 66 we also obtained spectra using the CD4 grating $\sim$ 6600–10600 )."1157 The observations of Be 20 and Be 29 were carried out in service mode during the period February-March 2006. while Cr 261 Mel 66 and were observed in May and December 2003.," The observations of Be 20 and Be 29 were carried out in service mode during the period February-March 2006, while Cr 261 Mel 66 and were observed in May and December 2003."1158 A log of observations (date. UT. exposure time. grating configuration. number of stars) is given in Table 2..," A log of observations (date, UT, exposure time, grating configuration, number of stars) is given in Table \ref{obslog}."1159 The spectra were reduced using the dedicated pipeline. and we analyzed the 1-d. wavelength-calibrated spectra using standard packages.," The spectra were reduced using the dedicated pipeline, and we analyzed the 1-d, wavelength-calibrated spectra using standard packages."1160 Radial velocities (RV) were derived with the task RVIDLINES in IRAF: corrections for the contribution of telluric lines were performed using TELLURIC and individual exposures were summed (see Paper and for more details)., Radial velocities $RV$ ) were derived with the task RVIDLINES in IRAF; corrections for the contribution of telluric lines were performed using TELLURIC and individual exposures were summed (see Paper and for more details).1161 To maintain the highest homogeneity in our work we tried to select always the same kind of stars. r.e.. red clump stars.," To maintain the highest homogeneity in our work we tried to select always the same kind of stars, i.e., red clump stars."1162 For the present observations we selected seven stars at the red clump of Mel 66. five stars at the red clump and one star near the RGB tip in Be 29. six stars at the clump and one slightly brighter 1n Cr 261. and six stars on the RGB near the clump level for Be 20 (since this cluster has no well defined clump).," For the present observations we selected seven stars at the red clump of Mel 66, five stars at the red clump and one star near the RGB tip in Be 29, six stars at the clump and one slightly brighter in Cr 261, and six stars on the RGB near the clump level for Be 20 (since this cluster has no well defined clump)."1163 Figures | and 2 show examples the spectra of all the stars observed. in a small wavelength region.," Figures 1 and 2 show examples the spectra of all the stars observed, in a small wavelength region."1164 We report information on the stars of the four clusters in Table 3:: we give the identification. equatorial coordinates (J2000). magnitudes (Cols.," We report information on the stars of the four clusters in Table \ref{dataobs}: we give the identification, equatorial coordinates (J2000), magnitudes (Cols."1165" 4-7). number of exposures for each star. the average heliocentric radial velocity (RV) and its rms. the S/N measured in the two wavelength regions around ~ 5600 and 6300Α.. and a note (M stands for ""member"". NM for ""non member"". NM?"," 4–7), number of exposures for each star, the average heliocentric radial velocity $RV$ ) and its rms, the $S/N$ measured in the two wavelength regions around $\sim$ 5600 and 6300, and a note (M stands for “member"", NM for “non member"", NM?"1166" for ""doubtful member"")."," for “doubtful member"")."1167 The B. V. and 7 magnitudes come from the different sources referenced below for each cluster. the K magnitudes are taken from 2MASS (Skrutskie et al. 2006)).7," The $B$ , $V$, and $I$ magnitudes come from the different sources referenced below for each cluster, the $K$ magnitudes are taken from 2MASS (Skrutskie et al. \cite{skrutskie})"1168 For Be 20 we adopted the identification number used for the FLAMES pointings. based on the EIS catalogue used to select and point stars.," For Be 20 we adopted the identification number used for the FLAMES pointings, based on the EIS catalogue used to select and point stars."1169 Only two stars of the six observed turned out to be secure cluster members., Only two stars of the six observed turned out to be secure cluster members.1170 These stars. 1201 and 1240. have RV of about 78.5 km s. that agrees with the average radial velocity found by Yong et al. (2005:," These stars, 1201 and 1240, have $RV$ of about 78.5 km $^{-1}$, that agrees with the average radial velocity found by Yong et al. \cite{yong05}:"1171 978.9. σΞ0.7 km s! ). therefore we label the two objects as cluster members: furthermore. 1201 and 1240 are the two stars closest to the cluster center.," +78.9, $\sigma$ =0.7 km $^{-1}$ ), therefore we label the two objects as cluster members; furthermore, 1201 and 1240 are the two stars closest to the cluster center."1172 Star 1240 has also been observed by Yong et al. (, Star 1240 has also been observed by Yong et al. (1173their star 8. as seen from their table 2).,"their star 8, as seen from their table 2)."1174 Stars 1401 and 1716 are clearly RV non members: the remaining objects. 1519 anc 1666. have RV differing by about 6-7 km s! from those of members.," Stars 1401 and 1716 are clearly $RV$ non members; the remaining objects, 1519 and 1666, have $RV$ differing by about 6-7 km $^{-1}$ from those of members."1175 Given the spectral resolution of our data. we car derive RV with a precision of better than 1 km κ. SO We conclude that 1519 and 1666 are most probably non members of Be 20.," Given the spectral resolution of our data, we can derive $RV$ with a precision of better than 1 km $^{-1}$, so we conclude that 1519 and 1666 are most probably non members of Be 20."1176 Another possibility is that they are long period binaries. but our data were obtained on too short a baseline to verify this possibility.," Another possibility is that they are long period binaries, but our data were obtained on too short a baseline to verify this possibility."1177 Magnitudes in the Johnson-Cousins system were taken from our reference photometry (MacMint et al., Magnitudes in the Johnson-Cousins system were taken from our reference photometry (MacMinn et al.1178 1994) for the two member stars., 1994) for the two member stars.1179 For stars 1666 and 1716 we used the V/c values from Andreuzzi et al. (, For stars 1666 and 1716 we used the $VI_{\rm C}$ values from Andreuzzi et al. (11802008). while for the remaining two objects we give in Table 3. the EIS magnitude. shifted to the MaeMinn et al.,"2008), while for the remaining two objects we give in Table \ref{dataobs} the EIS magnitude, shifted to the MacMinn et al."1181 and Kassis et al., and Kassis et al.1182 systems. respectively.," systems, respectively."1183 For Be 29 we selected five secure cluster members. às deduced from the RVs measured by Bragaglia et al. (," For Be 29 we selected five secure cluster members, as deduced from the $RV$ s measured by Bragaglia et al. ("11842005) and Carraro et al. (,2005) and Carraro et al. (11852004: our star 398. their 801).,"2004: our star 398, their 801)."1186 Also the sixth target (star 602). chosen among the red clump ones. turned out to be a member. às shown by the very similar RV.," Also the sixth target (star 602), chosen among the red clump ones, turned out to be a member, as shown by the very similar $RV$."1187 The average RV is 24.6640.37 km s!, The average $RV$ is $\pm$ 0.37 km $^{-1}$.1188 The identifications and BVc magnitudes in Table 3 are taken from the photometry by Tosi et al. (2004)).," The identifications and $BVI_{\rm C}$ magnitudes in Table \ref{dataobs} are taken from the photometry by Tosi et al. \cite{tosi04}) ),"1189 which is also in perfect agreement with Kaluzny (1994))., which is also in perfect agreement with Kaluzny \cite{kalBe29}) ).1190 The stars observed in Cr 261 share very similar RVs and therefore all of them are considered members (the average RV is —25.4341.11 km s7!)., The stars observed in Cr 261 share very similar $RV$ s and therefore all of them are considered members (the average $RV$ is $-$ $\pm$ 1.11 km $^{-1}$ ).1191 The ID adopted is the provisional one used in the FLAMES pointings. coming from the catalogue based on WFI@2.2m data.," The ID adopted is the provisional one used in the FLAMES pointings, coming from the catalogue based on WFI2.2m data."1192 Estimates of the reddening of Cr 261 are very uncertain. with values spanning from E(B— V)=0.22 up to 0.34 (see Spand 2005)).," Estimates of the reddening of Cr 261 are very uncertain, with values spanning from $E(B-V)$ =0.22 up to 0.34 (see Spanò \cite{tesi_spano}) )."1193 Recent studies (e.g. Carretta et al. 2005)), Recent studies (e.g. Carretta et al. \cite{carretta05}) )1194 seem to indicate a value close to 0.30. at least in the central region.," seem to indicate a value close to 0.30, at least in the central region."1195 We adopt here the WFI BV/ photometry calibrated by L. Prisinzano (private communication)., We adopt here the WFI $BVI$ photometry calibrated by L. Prisinzano (private communication).1196 As for as Mel 66. five of the observed stars have RV around 2] km s! (the mean RV is +21.2540.37 km s! y star 1865 has RV=+17.85 km s! and is therefore labeled as a doubtful member.," As for as Mel 66, five of the observed stars have $RV$ around 21 km $^{-1}$ (the mean $RV$ is $\pm$ 0.37 km $^{-1}$ ); star 1865 has $RV$ =+17.85 km $^{-1}$ and is therefore labeled as a doubtful member."1197 Star 1614 rotates very rapidly and was discarded from further analysis., Star 1614 rotates very rapidly and was discarded from further analysis.1198 For this cluster we adopted the V/c photometry by Kassis et al. (1997))., For this cluster we adopted the $VI_{\rm C}$ photometry by Kassis et al. \cite{kassis}) ).1199 The method of analysis is deseribed in Papers andmn: therefore we provide here only a brief summary. and we refer the reader to those papers for more complete information on the line lists adopted. references for atomic parameters and damping.," The method of analysis is described in Papers and; therefore we provide here only a brief summary, and we refer the reader to those papers for more complete information on the line lists adopted, references for atomic parameters and damping."1200 First. we determined thesolar abundances of Fe and other elements. in order to fix the zero points of the abundance scale and to minimize errors in the results.," First, we determined thesolar abundances of Fe and other elements, in order to fix the zero points of the abundance scale and to minimize errors in the results."

Showing the first 1,200 of 10358 lines. Download the file for the rest.